diff --git "a/0112.json" "b/0112.json" new file mode 100644--- /dev/null +++ "b/0112.json" @@ -0,0 +1,2602 @@ +{ + "0112/hep-ph0112253_arXiv.txt": { + "abstract": "s{ After brief personal recollections of the author's long-time friendship with Misha Marinov the problem of particle production by classical time- varying scalar field is discussed. In the quasiclassical limit the calculations are done by imaginary time method developed, in particular, in Marinov's works. The method permits to obtain simple analytical expressions which well agree with the later found numerical solutions. The results are compared with perturbative calculations and it is argued that perturbation theory gives an upper limit for the rate of production. } \\vspace{0.4cm} \\tableofcontents \\newpage ", + "introduction": " ", + "conclusions": "} It is demonstrated that imaginary time method very well describes particle production by scalar field. It is very simple technically and permits to obtain physically transparent results. The calculations here were done for a particular case of periodic or quasiperiodic oscillations of the field but, as shows the experience with production of $e^+e^-$-pairs by electric field (for a review see e.g. third paper in ref.\\cite{marinov72}), the method also works well in the opposite case of short pulse fields. The method is applicable in the quasiclassical limit. In the opposite case perturbation theory is applicable and hence one can obtain simple and accurate (semi)analytical estimates practically in all parameter range. The results of calculations in the quasiclassical limit are in a good agreement with subsequent numerical ones\\,\\cite{baacke98,green99}. An important difference between the latter papers and the initial one\\,\\cite{dolgov90} lays in the interpretation of the results. According to all these papers the occupation numbers of the produced particles quickly approaches unity but, in contrast to refs.\\cite{baacke98,green99}, it is argued in the paper\\,\\cite{dolgov90} that the total production rate is nevertheless suppressed in comparison to perturbation theory and the production of fermions by the inflaton with Yukawa coupling to fermions is always weak. This conclusion is verified above. As is shown in this paper, the occupation numbers may quickly reach unity both in perturbation theory and in non-perturbative case. Still even the production rate of particles obeying Boltzmann statistics is very weak to ensure fast (pre,re)heating. In the case of fermion production the rate is evidently much weaker because the production must stop when the occupation number reaches unity and to continue the process the produced fermions should be eliminated from the band. As is argued in sec.~\\ref{ss:small}, the non-perturbative effects can only diminish the production rate. The bosonic case is opposite: more bosons are in the final state, the faster is production. Thus even in perturbation regime the boson production can be strongly amplified because their occupation number may reach unity in much shorter time than $1/\\Gamma$ and the energy may be transferred from the inflaton to the produced bosons much faster than is given by the original perturbative estimates\\,\\cite{dolgov82}, where the effect of stimulated emission was not taken into account. Of course to realize this regime the band should not be destroyed by expansion and scattering, as argued in ref.~\\cite{dolgov90}. To summarize, we have shown that perturbation theory gives a good estimate of production of light fermions and bosons if Fermi exclusion principle or stimulated emission respectively are taken into account. The formally calculated production rate in perturbation theory is always larger than the non-perturbative one, at least in the simple cases that we have considered. So the results of perturbation theory may be used as upper bounds for production rates. Moreover, perturbation theory helps to understand physical meaning of the obtained results and to interpret them correctly. In many realistic cases (e.g. for large $g\\phi_0$ or $m_0$) perturbation theory is not applicable and to calculate the real production rate (not just an upper bound) one has to make more involved non-perturbative calculations. In quasiclassical (anti-perturbative) limit imaginary time method permits to obtain accurate and simple results and to avoid complicated numerical procedure" + }, + "0112/astro-ph0112416_arXiv.txt": { + "abstract": "We image 19 quasars with 22 \\dla \\ (DLA) systems using the F160W filter and the Near-Infrared Camera and Multiobject Spectrograph aboard the {\\it Hubble Space Telescope}, in both direct and coronagraphic modes. We reach 5$\\sigma$ detection limits of \\ab H=22 in the majority of our images. We compare our observations to the observed Lyman-break population of high-redshift galaxies, as well as Bruzual \\& Charlot evolutionary models of present-day galaxies redshifted to the distances of the absorption systems. We predict H magnitudes for our DLAs, assuming they are producing stars like an L$_{*}$ Lyman-break galaxy (LBG) at their redshift. Comparing these predictions to our sensitivity, we find that we should be able to detect a galaxy around 0.5-1.0 L$_{*}$(LBG) for most of our observations. We find only one new possible candidate, that near LBQS0010-0012. This scarcity of candidates leads us to the conclusion that most \\dla \\ systems are not drawn from a normal LBG luminosity function nor a local galaxy luminosity function placed at these high redshifts. ", + "introduction": "In recent years, our knowledge of the high redshift universe has finally grown to include a significant number of galaxies powered mainly by star formation, often referred to as Lyman-break galaxies because most were discovered by searching for the Lyman-limit discontinuity using broad band colors \\citep{ste96}. There is still the concern, however, that these galaxies may not be representative of the typical galactic mass at that epoch. Dust, for example, could strongly affect the discovery rate of galaxies found using searches, like the Lyman-break method, that search in rest frame ultraviolet light \\citep{row97}. Quasar absorption line studies, on the other hand, have the advantage of being able to follow the neutral hydrogen content of the universe regardless of whether it emits light or not. While absorption line studies may also suffer from selection effects, with dustier intervening gas creating dimmer quasars that are either less well studied or not discovered at all \\citep{pei95, car98}, they represent a potentially less biased tracer of the content of the universe than the Lyman-break galaxy. Examinations of \\dla \\ (DLA) systems, which make up the bulk of all absorption line neutral hydrogen, show that at high redshifts, the mass of neutral hydrogen seen in DLAs per unit comoving volume is roughly the same as the mass density seen in luminous matter (i.e. stars) in present-day spiral galaxies \\citep{lan91}. Since this is a much larger quantity of neutral gas than we observe in galaxies today \\citep{rao93, zwa98}, this suggests that much of the gas we see in high z DLAs has ended up in stars. In fact, \\cite{wol95} identify the DLA systems as the likely progenitors of current spiral galaxies. However, some recent work has challenged this idea, showing no drop in the cosmological mass density of neutral hydrogen from redshifts of 3 down to 0.5 \\citep{rao00}, instead of the expected gradual decline seen in other work \\citep{smi96,jan98,sto00}. Another indication that DLAs may not be forming today's galaxies is the lack of any metal enrichment in lower redshift DLAs, with the mean Zn/H ratio remaining around 1/10 solar from z=1-3 \\citep{pet99}. It should be noted that the statistical uncertainties involved in these measurements remain significant, since all work on DLAs at low redshifts relies on a small number of objects (23 with z$<$1.65 for \\cite{rao00}), with the handful with highest column densities dominating the statistics. These low redshift DLA studies are also the most vulnerable to bias from any possible obscuration of quasars by dust in foreground absorbers \\citep{sto00, pei95}. Modeling of velocity profiles has not settled the question. Supporting the contention that DLAs form spiral galaxies, \\cite{pro97} demonstrate how models of rotating disks, not unlike today's spirals, fit the observed double trough velocity profiles seen in some DLA metal lines. However, there are many other viable models that also produce similar velocity profiles, such as merging protogalactic clumps \\citep{hae98}, randomly moving clouds in spherical halos \\citep{mcd99}, and multiple gaseous discs in a single halo \\citep{mal00}. Another problem with the DLAs as spiral galaxies hypothesis is the number of DLA systems that are being discovered, compared with the number of spiral galaxies around at the present day. Based on the present day density and size of spiral galaxies, one can estimate the number that should intercept a random line of sight to a quasar at some high redshift. While the exact prediction depends on the assumed cosmology, \\cite{lan91} found that the number of DLA systems being discovered is at least twice that predicted, even under the more favorable cosmological assumptions. This means that if DLAs are high-redshift spiral galaxies, then there must be strong evolution in the spiral population with redshift. Either spiral galaxies were more numerous in the past, or they were larger or both. Another way around this discrepancy is to assume that our understanding of local galaxies and the measured fiducial parameters used for these predictions are flawed. For instance, a large population of low surface brightness galaxies \\citep{imp89} or low luminosity dwarf galaxies \\citep{cen98} could dominate the absorption cross sections responsible for \\dla \\ absorption lines. This, however, is not the case for low-redshift Mg II absorption line systems for which visible counterparts have been found, where the majority are essentially bright galaxies with typical luminosities just under L$_{*}$ \\citep{ber91}. It would be valuable to find some connection between DLAs and high-redshift galaxies, our two largest sources of information on the non-AGN high-redshift universe. Unfortunately, while we know a great deal about the distribution, metallicity, and evolution of the neutral gas in high-redshift DLAs, we know almost nothing about their associated stars, galactic sizes, or morphologies. Only by imaging starlight from these DLA systems can we measure how much, if any, star formation is actually occurring in these objects. Lyman break galaxies produce stars at a steady pace of about 10 M$_{\\sun }$ yr$^{-1}$, depending heavily on cosmology and the large dust extinctions assumed \\citep{see96}, not that different from the \\ab 10 M$_{\\sun }$ yr$^{-1}$ seen in H$\\alpha$ emission line galaxies today \\citep{gal95}. If DLAs are progenitors of present day galaxies and/or the same population as the Lyman break galaxies, then they should also be producing stars and starlight. This is true whether the DLAs are galactic disks, spheroids, or some other assembled structure. Searches for visible candidates for low to medium-redshift DLA systems have had mixed results. Often candidates are found, but less than half of them are suspected to be reasonably bright spiral galaxies. The rest are a non-uniform collection of low surface brightness galaxies, bright compact objects, and dim dwarfs \\citep{ste94,leb97,lan97,rao98,pet00}. Again the statistics are poor (less than 15 objects) and most have not been spectroscopically confirmed, leaving the nature of these DLAs still murky. High-redshift systems have been even more problematic, with much lower candidate discovery rates. A major reason is that this is a difficult measurement, considering the large distances to these objects and their close projected proximity to a bright quasar. A common search technique has been to look for emission lines, Lyman \\a \\ or H\\a , where the emission line flux ought to stand out from the background and the quasar light may be depressed. These searches have found at least two fairly unambiguous detections, near quasars PKS0528-250 \\citep{mol93} and 2231+131 \\citep{djo96}, but broader searches of DLA samples have generally found nothing at all \\citep{low95,bun99}. Another strategy for imaging DLAs has been to go into the infrared, where K-corrections would be more favorable and dust, often cited as a possible explanation for the failure of Lyman \\a \\ searches, would likely be only a minor effect. For instance, \\cite{ara96} observed ten quasars with DLA systems in the near-infrared, discovering two candidates (in front of quasars 0841+129 and 1215+333) after subtracting off the PSF of quasar. Since both candidates lie about an arcsecond away from the quasar, the PSF subtraction is critical and an error could lead to a false detection. Higher resolution imaging would greatly help. We have obtained imaging of 19 quasars with confirmed DLA systems using the Near Infrared Camera and Multi-Object Spectrometer (NICMOS) onboard the Hubble Space Telescope (HST). This gives us the advantage of both the near-infrared, with its favorable K-corrections and low dust dependence, and high spatial resolution, which will separate any candidates from the quasar's light, even at distances of less than arcsecond. In addition, the majority of the observations were done using the NICMOS coronagraph, which greatly decreases scattered light from the quasar, improving chances of candidate detection. If DLAs are high-redshift disk galaxies producing stars, missed only because of their close proximity to the quasar, a survey of this type can uncover them. Another advantage of a large survey of a sample like ours is that it presents the failures as well as the successes, something that is often missing from papers on individual candidate discoveries. A rate of DLA system discoveries gives us a statistically significant result for the entire DLA class of objects, something an individual candidate can not. Section 2 describes our sample and the observations, while $\\S$3 describes the data reduction and analysis. Section 4 discusses the sensitivity of our survey to these hypothetical L$_{*}$(LBG) galaxies. Section 5 presents the results of our candidate search. Finally, $\\S$ 6 presents our discussion and conclusions. (A list of notes on individual DLAs is provided in the Appendix). ", + "conclusions": "We are not seeing galaxies near quasar lines of sight that contain \\dla \\ absorption. This result is very similar to the negative results of groups surveying DLA systems searching for the emission lines of Lyman \\a \\ \\citep{low95} and H\\a \\ \\citep{mal95,bun99}, suggesting that that their difficulties may not have been because the DLA systems do not produce strong emission lines, but because they do not produce much light at all. Our ability to look much closer to the quasar, 0.5 arcseconds or better, also weakens the argument that DLAs are not being found because they lie at such small impact parameters. In fact, the candidate seen by \\cite{ara96} near Q1215+333 at a distance of 1.3 arcseconds is not seen in our data. With a K magnitude of 20.1, we should have picked it up, even with a red color (H-K $\\geq$ 1.5). We tried models creating all stars in a single burst around z=10, but none were redder than H-K = 1.1 at z=2 \\citep{bru96}, with most of our exponential models predicting H-K = 0.8-0.9. Colors redder than H-K = 1.2 required redshifts of z=3-3.5 and produced correspondingly dimmer galaxies. This makes it likely the former detection was a result of the very difficult subtraction of a bright PSF so near the quasar. It does not seem likely that we can continue to hide the majority of DLA systems by pushing them all into the tiny area (less than 0.8 arcsec$^2$ for our cornagraphic images) directly in front of the quasar, especially considering the extreme number evolution of galaxies that would require. Also not seen are the \\cite{ste92} candidate for [HB89] 0000-263 (DLA \\# 1), at a distance of 2.8$\\arcsec$, and the \\cite{ste95} and \\cite{leb97} candidate for [HB89] 0454+039 (DLA \\# 6), measured at two different distances from the quasar, 2.1$\\arcsec$ and 0.8$\\arcsec$ respectively. These last two may have been missed if the candidate were blue enough, R-H $\\leq$ 3.4 for the [HB89] 0000-263 candidate and R-H $\\leq$ 2.3 for [HB89] 0454+039. Models of high-redshift colors \\citep{bru96} suggest that the former is likely (predicted R-H = 1.3-1.5 for exponential star formation models at z=3.4), while the latter is possible, but difficult to fit (predicted R-H = 2.3-3.5 for exponential star formation models at z=0.86). Galaxies as blue as R-H = 2.3 at z=0.86 require either strong continuous star formation or a strong starburst just prior to observation. All this leads us to speculate as to why we are not detecting light from almost any of these DLA systems. One possibility is that they not producing much light because they have not started making many stars by the period at which we are observing them. Most of our DLA sample comes from around z$\\sim$2, so if most star formation occurred after that we would not be able to see it. One possible problem with this scenario is the previously observed decrease in comoving mass density of neutral gas seen in DLAs from z$\\sim$3.5 to z=2 and then down into the present day. If DLAs are spiral galaxy progenitors, this is explained as the transformation of their neutral gas into stars. Initial work by \\cite{wol95} indicated that as much one half of the neutral gas was depleted before z=2, meaning half of all stars should have been created before then and would therefore likely be visible to our survey. However, further work by \\cite{smi96}, using some of the highest redshift DLAs known, showed a less dramatic neutral gas evolution, with only 20\\% of the neutral gas processed by z=2. Also, as noted in the introduction, \\cite{rao00} puts the whole idea of the decline of neutral gas in DLAs with time into doubt, showing no decline down into low redshifts. Neutral gas evolution aside, if DLAs are to become today's galaxies they must form stars at some point and a large amount of star formation is already measured to be underway before z=2 \\citep{mad98}. Waiting until low redshifts to produce the vast majority of a galaxy's stars would create a problem with the number of highly luminous galaxies required at z$<$2. Dust could also play a factor, absorbing light even at rest wavelength optical wavelengths. However there is no evidence for such large quantities of dust seen in the quasar spectra \\citep{pei91,pet97}, requiring a most unlikely conspiracy of much higher dust quantities surrounding the star formation regions, but not along the quasar line of sight. Another possibility is that the DLA systems are not compact, high surface brightness objects like the spiral galaxies we study, but are instead diffuse, low-surface-brightness galaxies as suggested by \\cite{jim99}. These galaxies would be inefficient at transforming their gas into stars, so would take longer to start forming stars, and would be harder to detect, even if their total magnitude equalled our predictions. If this is the case, the majority of DLAs are not like the LBGs discovered so far, and either these low surface brightness galaxies have to evolve into the high surface brightness spirals we see around us today, or these DLA systems will never become today's L$_{*}$ galaxies. This possibilty -- that DLAs are not high-redshift versions of present day disk galaxies -- could by itself explain all of our upper limits. They might be protogalactic blobs that will eventually fall into galaxies, helping them form or grow, or they might never become part of standard galaxies at all, being something entirely separate. While there may be a handful of DLA systems that we can detect in emission, the vast majority do not seem to be taken from the distributions of either $z=0$ or $z=3$ galaxies. Either they are not producing as many stars, something is absorbing their light, or their surface brightness is unusually low. From the evidence gathered so far, it appears that the distribution of DLA systems is inconsistent with both the evolution of Lyman-break galaxies forward in time and present day galaxies backwards." + }, + "0112/astro-ph0112550_arXiv.txt": { + "abstract": "An overview of the current status of WIMP direct searches is presented, emphasizing strategies, achievements and prospects. ", + "introduction": "Experimental observations and robust theoretical arguments have established that our universe is essentially non-visible, the luminous matter scarcely accounting for one per cent of the energy density of a flat universe. The distribution of a flat universe ($\\Omega=\\Omega_{M}+\\Omega_{\\Lambda}=1$) attributes to the dark energy about $\\Omega_{\\Lambda}\\sim70\\%$, whereas the matter density takes the remaining $\\Omega_{M}\\sim30\\%$, consisting of, both, visible ($\\Omega_{l}\\sim0.5\\%-1\\%$) and non-visible (dark) matter. This dark component consists of ordinary baryonic matter ($\\Omega_{B}\\sim 4-5\\%$), (possibly made by machos, jupiters, dust, black holes, etc.) and a large fraction (up to $\\Omega_{NB}\\sim25\\%$) of non baryonic dark matter, supposedly made by non-conventional, exotic particles. The minimal requirements to be fulfilled by the non-baryonic dark particles are to provide the right relic abundance, to have non-zero mass, zero electric charge and a weak interaction with ordinary matter. There are several candidates to such species of matter provided by schemes beyond the Standard Model of Particle Physics. The galactic DM axions, the SUSY WIMPs (like neutralinos) and the light neutrinos (of some non standard models) are particularly attractive. The WIMPs, Weak Interacting Massive Particles, are favorite of experimentalists ($\\sim20$ experiments for WIMPs vs. 2 for axions) and so of this talk. This talk deal with the recent efforts done in its direct search illustrated by a few, selected experiments. \\begin{table*}[htb] \\caption{WIMP Direct Detection in underground facilities experiments currently running (or in preparation) } \\label{table:2} \\begin{tabular}{lll} \\hline LABORATORY & EXPERIMENT & TECHNIQUE \\\\ \\hline BAKSAN (Russia) & IGEX & 3$\\times$1 Kg Ge-ionization \\\\ BERN(Switzerland) & ORPHEUS &(SSD) Superconducting Superheated Detector, 0.45 Kg Tin \\\\ BOULBY & NaI & NaI scintillators of few Kg (recently completed) \\\\ (UK)& NAIAD & NaI unencapsulated scintillators (50 Kg)\\\\ & ZEPLIN & Liquid-Gas Xe scintillation/ionization I: 4 Kg single phase \\\\ & & II: 30 Kg Two phases \\\\ & DRIFT & Low pressure Xe TPC (in preparation) 1 m$^{3}\\rightarrow$ 10 m$^{3}$ \\\\ CANFRANC & COSME & 234 g Ge ionization \\\\ (Spain) & IGEX & 2.1 Kg Ge ionization \\\\ & ANAIS & 10$\\times$10.7 Kg NaI scintillators \\\\ & ROSEBUD & 50g Al$_{2}$O$_{3}$ and 67g Ge thermal detectors \\\\ & & CaWO$_{4}$ 54g and BGO 46g scintillating bolometers \\\\ FREJUS/MODANE & SACLAY-NaI & 9.7 Kg NaI scintillator (recently completed) \\\\ (France) & EDELWEISS I &70 g Ge thermal+ionization detector \\\\ & EDELWEISS II & 4$\\times$320 g Ge thermal+ionization detectors \\\\ GRAN SASSO &H/M &2.7 Kg Ge ionization \\\\ (Italy) &HDMS & 200g Ge ionization in Ge well \\\\ & GENIUS-TF & 40$\\times$2.5 Kg unencapsulated Ge (in preparation) \\\\ & DAMA & NaI scintillators (87.3 Kg) \\\\ & LIBRA & NaI scintillators 250 Kg (in preparation) \\\\ & Liquid-Xe & Liquid Xe scintillator (6 Kg) \\\\ & CaF$_{2}$ & Scintillator \\\\ & CRESST I & (4$\\times$260g) Al$_{2}$O$_{3}$ thermal detectors \\\\ & CRESST II & Set of 300g CaWO$_{4}$ scintillating bolometers (up to 10 Kg) \\\\ & MIBETA & 20$\\times$340g TeO$_{2}$ thermal detector \\\\ & CUORICINO & 56$\\times$760g TeO$_{2}$ thermal detector (being mounted) \\\\ & CUORE &1000$\\times$760g TeO$_{2}$ (in preparation) \\\\ RUSTREL (France)& SIMPLE & (SDD)Superheated Droplets Detectors (Freon) \\\\ STANFORD UF/ & CDMS - I & 100g Si; 6$\\times$165g Ge thermal+ionization detectors \\\\ SOUDAN(USA)& CDMS - II & 3$\\times$250g Ge and 3$\\times$100g Si Thermal+Ionization \\\\ SNO (Canada) & PICASSO &(SDD)Superheated Droplets Detectors (1.34g of Freon) \\\\ OTO & ELEGANTS-V & Large set of massive NaI scintillators \\\\ (Japan) &ELEGANTS-VI & CaF$_{2}$ scintillators \\\\ \\hline \\end{tabular}\\\\[2pt] \\end{table*} Galactic halo WIMPs could be directly detected by measuring the nuclear recoil produced by their elastic scattering off target nuclei in suitable detectors at a rate which depends of the type of WIMP and interaction. In the case of WIMPs of $m\\sim GeV~\\mathrm{to}~\\textit{TeV}$ and $v\\sim10^{-3}c$ the nuclear recoil in the laboratory frame $E_{R}=\\frac{\\mu^{2}}{M}v^{2}(1-cos\\theta)$ is in the range from 1 to 100 \\textit{KeV}. $M$ is the nuclear mass, $\\mu$ the ($m, M$) reduced mass and $\\theta$ the WIMP-nucleus (c. of m.) scattering angle. Only a fraction $QE_{R}=E_{vis} (\\equiv E_{eee})$ of the recoil energy is visible in the detector, depending on the type of detector and target and on the mechanism of energy deposition. The so-called Quenching Factor Q is essentially unit in thermal detectors whereas for the nuclei used in conventional detectors it ranges from about 0.1 to 0.6. The energy delivered by the WIMP results in a small signal (1-100 KeV) which shows up even smaller (for Q$<$1). Moreover this signal falls in the low energy region of the spectrum, where the radioactive and environmental background accumulate at much faster rate and with similar shape. That makes WIMP signal and background practically undistinguishable. On the other hand, the smallness of the neutralino-matter interaction cross-section implies that the process looked for is very rare. Customarily, one compares the predicted event rate with the observed spectrum. If the former turns out to be larger than the measured one, the particle which would produce such event rate can be ruled out as a Dark Matter candidate. That is expressed as a contour line $\\sigma$(m) in the plane of the WIMP-nucleon elastic scattering cross section versus the WIMP mass. That excludes, for each mass m, those particles with a cross-section above the contour line $\\sigma$(m). The level of background sets, consequently, the sensitivity of the experiment in eliminating candidates or in constraining their masses and cross sections. However, this simple comparison will not be able to identify the WIMP. A convincing proof of its detection should be provided by a distinctive signature characteristic of WIMPs. Such distinctive labels do exist: they are originated by the motion of the Earth in the galactic halo \\cite{Drukier,Spergel}. These signatures are an annual modulation of the rate and a directional asymmetry of the nuclear recoil. Narrowing first the window of the possible WIMP existence and looking then for its identification is the purpose of the experimental searches. Table \\ref{table:2} gives an overview of the experiments on direct detection of WIMPs currently in operation or in preparation, which is the subject of this talk. General reviews for WIMP dark matter are given in Ref. \\cite{Gri}. WIMP direct detection is reviewed, for instance in Ref. \\cite{Mor2,Mor}. WIMPs can be also looked for, indirectly, in the galactic halo, looking for its presence in cosmic ray experiments in terms of antiprotons, positrons or gamma rays produced by WIMP annihilation in the halo. One can also search in underground, underwater or under-ice detectors, looking (also indirectly) for WIMPs through the high energy neutrinos emerging as final products of the WIMP annihilation in celestial bodies (Earth or Sun). ", + "conclusions": "The direct search for WIMP dark matter proceeds at full strength. More than twenty experiments on direct detection illustrate the effort currently being done. New, dedicated experiments are focusing now in the identification of WIMPs, discriminating the nuclear recoils from the background, rather that in constraining or excluding their parameters space. Their current achievements and the projections of some of them have been reviewed in this talk. \\begin{figure}[b] \\centerline{\\includegraphics[height=4.5cm]{fig13.eps}} \\caption{} \\label{fig13} \\end{figure} The present experimental situation can be summarized as follows: the rates predicted for SUSY-WIMPs extend from 1-10 c/Kgday down to $10^{-4}-10^{-5}$ c/Kgday, in scatter plots, obtained within MSSM as basic frame implemented in various alternative schemes. A small fraction of this window is testable by some of the leading experiment. The rates experimentally achieved stand around 1 c/Kgday (0.1 c/Kgday at hand) (CDMS, EDELWEISS) and differential rates $\\sim0.1-0.05$ \\textit{c/KeV Kg day} have been obtained by IGEX and H/M, in the relevant low energy regions. The deepest region of the exclusion plots achieved stands around a few $\\times10^{-6}$pb, for masses 50-200 GeV (DAMA, CDMS, EDELWEISS, IGEX). The current status of the best exclusion plots is depicted comparatively in Fig. \\ref{fig13}. There exists an unequivocal annual modulation effect (see Fig. \\ref{fig4}) reported by DAMA (four yearly periods), which has been shown to the compatible (DAMA) with a neutralino-WIMP, of m$\\sim50-60$GeV and $\\sigma^{Si}_{n}\\sim 7\\times10^{-6}$pb. Recent experiments exclude at greater or lesser extend (CDMS, EDELWEISS, IGEX) the DAMA region. To reach the lowest rates predicted ($10^{-5}$ c/Kgday) in SUSY-WIMP-nucleus interaction, or in other words, to explore coherent interaction cross-sections of the order of $10^{-9}-10^{-10}$pb, substantial improvements have to be accomplished in pursuing at its best the strategies reviewed in this talk, with special emphasis in discriminating the type of events. These strategies must be focussed in getting a much lower background (intrinsic, environmental, ...) by improving radiopurity and shieldings. The nuclear recoil discrimination efficiency should be optimized going from above 99.7\\% up to 99.9\\% at the same time that the energy $E_{vis}$ at which discrimination applies should be lowered. The measurement of the parameters used to discriminate background from nuclear recoils should be improved and finally one needs to increase the target masses and guaranty a superb stability over large exposures. With these purposes various experiments and a large R+D activity are under way. Some examples are given in Table \\ref{table:1}. The conclusion is that the search for WIMPs is well focused and should be further pursued in the quest for their identification. \\begin{table*}[htb] \\caption{WIMP Direct Detection Prospect} \\label{table:1} \\begin{tabular}{ll} \\hline &BEING INSTALLED/OR PHASE II EXPERIMENTS~(To start 2001-2002) \\\\ \\hline CDMS-II & (Ge,Si) Phonons+Ioniz 7 Kg, B$\\sim 10^{-2}-10^{-3}$ c/Kgd, $\\sigma \\sim 10^{-8}$ pb \\\\ EDELWEISS-II & (Ge) Phonons+Ioniz 6.7 Kg, B$\\sim 10^{-2}-10^{-3}$ c/Kgd, $\\sigma \\sim 10^{-8}$ pb (40-200 GeV) \\\\ CUORICINO & TeO$_{2}$ Phonons 42 Kg, B$\\sim 10^{-2}$ dru, $\\sigma \\sim 0^{-7}$ pb \\\\ CRESST-II & CaWO$_{4}$ Phonons+light, $B<10^{-2}-10^{-3}$ dru (15 KeV), $\\sigma \\sim 10^{-7}-10^{-8}$ pb \\\\ & (50-150 GeV) \\\\ IGEX & Ge Ioniz 2.1 Kg, B$< 10^{-1}-10^{-2}$ dru, $\\sigma \\sim 2\\times 10^{-6}$ pb (40-200 GeV) \\\\ HDMS & Ge Ioniz 0.2 Kg, $\\sigma \\sim 6 \\times 10^{-6}$ pb (20-80 GeV) \\\\ ANAIS & NaI Scintillators 107-150 Kg, B(PSD)$\\leq 0.1$ dru, $\\sigma \\sim 2 \\times 10^{-6}$ pb \\\\ NAIAD & NaI Scintillators 10-50 Kg, B(PSD)$\\leq 0.1$ dru, $\\sigma \\sim 10^{-6}$ pb (60-200 GeV) \\\\ \\hline & IN PREPARATION~(To start 2002-2003) \\\\ \\hline LIBRA (DAMA) & NaI Scintillators 250 Kg \\\\ GENIUS-TF & Ge Ioiniz 40 Kg, B$<10^{-2}$ dru, $E_{Thr}$=10 KeV $\\rightarrow \\sigma \\sim 10^{-6}$ pb (40-200 GeV), \\\\ & $E_{Thr}$=2 KeV $\\rightarrow \\sigma \\sim 10^{-7} pb$ (20-80 GeV) \\\\ ZEPLIN-II & Xe-Two-phase 40 Kg, NR discrim$>$99\\%, B$<10^{-2}$ dru, $\\sigma \\sim 10^{-7}$ pb \\\\ DRIFT-I & Xe TPC 1 m$^{3}$, B$<10^{-2}$ dru, $\\sigma \\sim 10^{-6}$ pb (80-120 GeV) \\\\ \\hline & IN PROJECT~($>$2003-2005) \\\\ \\hline CUORE & TeO$_{2}$ Phonons 760 Kg, $E_{Thr}\\sim$2.5 KeV, B$\\sim 10^{-2}-10^{-3}$ dru, $\\sigma \\sim 5\\times 10^{-8}$ pb \\\\ GENIUS 100 & Ge ioniz 100 Kg, $E_{Thr}\\sim$10 KeV \\\\ (GENINO) & B$\\sim 10^{-3}-10^{-5}$ dru, $\\sigma \\sim 5\\times 10^{-8}-2\\times 10^{-9}$ pb \\\\ GEDEON & Ge ioniz 28-112 Kg, B$\\sim 2\\times 10^{-3}$ dru ($>$10 KeV) $\\sigma \\sim 10^{-7}-10^{-8}$ pb (40-200 GeV) \\\\ \\hline & THE FUTURE~($>$2005-2007) \\\\ \\hline DRIFT 10 & Xe 10 m$^{3}$ TPC, $\\sigma \\sim 10^{-}8$ pb \\\\ ZEPLIN-MAX & Xe Two-Phase, $\\sigma \\sim 10^{-10}$ pb \\\\ GENIUS & Ge ioniz 1-10 Tons, $\\sigma \\sim 10^{-9}-10^{-10}$ pb \\\\ DRIFT-1 ton & Xe 1 Ton TPC, $\\sigma \\sim 10^{-10}-10^{-11}$ pb \\\\ \\hline \\end{tabular}\\\\[2pt] \\end{table*}" + }, + "0112/astro-ph0112146_arXiv.txt": { + "abstract": "We present spectroscopic and high speed photometric data of the eclipsing polar V895 Cen. We find that the eclipsed component is consistent with it being the accretion regions on the white dwarf. This is in contrast to Stobie et al who concluded that the eclipsed component was not the white dwarf. Further, we find no evidence for an accretion disc in our data. From our Doppler tomography results, we find that the white dwarf has $M$\\gtae 0.7\\Msun. Our indirect imaging of the accretion stream suggests that the stream is brightest close to the white dwarf. When we observed V895 Cen in its highest accretion state, emission is concentrated along field lines leading to the upper pole. There is no evidence for enhanced emission at the magnetic coupling region. ", + "introduction": "Polars are interacting binaries consisting of a red dwarf secondary and a strongly magnetized ($\\sim$10-200 MG) white dwarf primary. In these systems the secondary fills its Roche lobe. Material falls under gravity from the secondary towards the primary, initially along the binary orbital plane before the magnetic field of the primary forces it to leave the orbital plane and eventually impacts quasi-radially onto the white dwarf. Unlike non-magnetic cataclysmic variables (CVs) the strong magnetic field of the white dwarf is high enough to prevent the formation of an accretion disc around the primary. Although our understanding of the accretion flow near the white dwarf is now relatively well understood (eg Wu 2000 and references therein) the region where accretion flow first interacts with the magnetic field of the white dwarf is not. The flow interacts with the magnetosphere in a complex manner and it is not easy to isolate stream emission from other emission sources in the system (which are generally brighter). Eclipsing systems provide an opportunity to study the accretion flow as a separate and distinct source for a short period of time, when the emission from the bright accretion region on the white dwarf is blocked by the secondary. Light curves of the eclipse contain information about the structure and the brightness distribution along the stream. The stream brightness distribution can be retrieved using an indirect imaging technique which can reconstruct the brightness of the region between the primary and the secondary. One such technique is that of Hakala (1995) who devised an indirect imaging method based on Maximum Entropy to deduce the brightness distribution along the accretion stream of HU Aqr. This technique has been developed further by Harrop-Allin et al (1999a, 1999b, 2001) who used a more physically realistic stream trajectory, improving the model's optimizing algorithm and including the projection effects. V895 Cen was discovered serendipitously using the {\\sl EUVE} satellite (Craig et al 1996). Craig et al observed strong line emission and low/high brightness states and concluded V895 Cen was likely to be a polar. Its orbital period of $P_{orb}$=4.765 h (Stobie et al 1996) is the second longest period currently known for a polar and the system is found to alternate frequently between high and low accretion states. Stobie et al (1996) concluded that it is an eclipsing polar but that the secondary minimum of the ellipsoidal variation was offset with respect to the eclipse. They suggested the eclipsed component was a hot compact source which appeared to be distinct from the white dwarf, probably associated with the accretion stream. Howell et al (1997) also concluded the eclipse to be of an extended object much larger than the white dwarf, perhaps a partial accretion disc which forms during the high state. Stobie et al (1996) found no evidence for significant levels of polarization in a low accretion state. To investigate the nature of the eclipsed component, we have applied the techniques of indirect imaging and Doppler tomography to this system to study the spatial and temporal changes in the stream and to determine if the eclipse is associated with the white dwarf or the stream. ", + "conclusions": "\\subsection{The eclipsed source} Stobie et al (1996) suggested that the eclipsed component was not the white dwarf. This was based on the fact that the eclipse occurred before the secondary minimum in the light curve. Assuming the secondary minimum was the true marker of inferior conjunction, and taking the timings of the eclipse ingress and egress relative to this, they concluded that the eclipsed source was $\\sim$30 white dwarf radii from the white dwarf. We have found that the phasing of the eclipse ingress and egress occurs at exactly the same orbital phase in both low, intermediate and high state data (cf Figure 5). If the eclipsed source was the stream as suggested by Stobie et al (1996) we would not expect to observe the stability in the eclipse features as we do. We now address the fact that the eclipse appears off-set from the secondary minimum in the low state light curves of Stobie et al (1996). Many polars show evidence for heating of the trailing face of the secondary by the accretion region on the white dwarf. It is expected that even if the irradiation is sharply reduced or switched off, the trailing face of the secondary will remain heated for some duration. Szkody et al (1999) estimate that in the case of the polar AR UMa it takes around 5 months for the secondary to cool down to the temperature of the unheated part of the star. Indeed, from our Doppler tomography results (\\S \\ref{tomo}) we find that the secondary in V895 Cen is still heated when we observed it in a low accretion state. This has the effect of increasing the optical flux between $\\phi$=0.5--0.9 compared to a secondary star with no irradiation. We suggest that the apparent offset between the eclipse and the secondary minimum is due to asymmetric irradiation of the secondary star. \\subsection{The indirect stream mapping results} Our results show that in the intermediate/low accretion state the stream brightens rapidly as it nears the white dwarf. In our brightest state data, we find that stream emission is concentrated mainly along the field lines leading to the upper pole. At face value this suggests that as the system reaches a high enough mass transfer rate, the accretion mode goes from a two-pole to a one-pole model. It is not clear why this would be the case, but it suggests that with an increase in the mass transfer rate, the upper pole is now the more favorable pole to accrete. In all our model fits, there is no evidence for a brightening of the accretion stream at the magneto-spheric interaction region. These results are similar to that of the low accretion state data of HU Aqr (Harrop-Allin et al.~ 2001). However, there was some indication that the stream brightened at the interaction region in the $U$ and $B$ bands. Using emission line data of HU Aqr in a high accretion state and a different technique to that used here, Vrielmann \\& Schwope (2001) derived stream brightness maps and found a brightening of the stream in the magneto-spheric interaction region of HU Aqr. It is possible that the accretion state (and hence amount of irradiation) plays an important role in determining whether the stream is found to brighten at the interaction region. It is interesting to compare the coupling radius that we derive from our model fits with that of HU Aqr (Harrop-Allin 1999b, 2000). They find that in a high accretion state, $R_{\\mu}\\sim$0.18a. We find a mean value of $R_{\\mu}\\sim$0.25 from our fits. HU Aqr has a magnetic field strength of 36 MG (Schwope, Thomas \\& Beuermann 1993). Since we expect the accretion flow to interact with the magnetic field when the magnetic pressure equals the ram pressure of the flow, we predict that the magnetic field strength of V895 Cen is significantly larger than that of HU Aqr, other things (such as \\Mdot) being equal." + }, + "0112/astro-ph0112364_arXiv.txt": { + "abstract": "Recent stellar population and chemical abundance studies point to an accreted origin of $\\omega$ Cen. In this light, and given the retrograde, small size orbit of $\\omega$ Cen, we search for a kinematical signature left by its hypothetical parent galaxy in the Solar neighborhood. We analyze the largest-to-date sample of metal poor stars (Beers {\\it et al.} 2000) and we find that, in the metallicity range $-2.0 <$ [Fe/H] $\\le -1.5$, a retrograde signature that departs from the characteristics of the inner halo, and that resembles $\\omega$ Cen's orbit, can be identified. ", + "introduction": "Recent advances in understanding the nature and origin of the highly unusual globular cluster $\\omega$ Centauri (see Majewski {\\it et al.} 2000 for a summary of properties), are due primarily to the following findings: the multiple-peak metallicity distribution seen in the structure of the giant branch (Lee {\\it et al.} 1999; Pancino {\\it et al.} 2000; Frinchaboy {\\it et al.} 2001), the correlation between age and metallicity (e. g., Hughes \\& Wallerstein 2000, Hilker \\& Richtler 2000), and the s-process enhanced enrichment in cluster stars compared to halo stars of similar metallicity (Smith {\\it et al.} 2000; Vanture, Wallerstein \\& Brown 1994). These findings suggest that $\\omega$ Cen underwent self-enrichment with at least three primary enrichment peaks (Pancino {\\it et al.} 2000, Frinchaboy {\\it et al.} 2001), over a period of at least 3 Gyr (Hughes \\& Wallerstein 2000). The s-process heavy-elements are primarily synthesized in low-mass (1.5 to 3.0 M$_{\\odot}$) asymptotic giant branch (AGB) stars (see e.g., Travaglio {\\it et al.} 1999 and references therein). In order to enrich the cluster in s-process elements, the ejecta from low-mass stars that evolve on timescales of $10^{9}$ years had to be retained by the cluster and incorporated in the next generations of stars. This long and complex star formation history is inconsistent with the cluster originating on its current orbit, which is of low energy and confined to the disk. With a period of only 120 Myr (Dinescu, Girard, \\& van Altena 1999 - hereafter DGvA), the frequent disk crossings would have certainly swept out all of the intracluster gas soon after its formation, and the result would be a single-metallicity system that would resemble most of the Galactic globular clusters. It appears thus that $\\omega$ Cen evolved somewhere away and independently from the Milky Way, in a system that was massive enough to retain ejecta from previous generations of stars, and to undergo multiple episodes of star formation. Its current orbit can be reconciled with the complex star formation history only if it represents a strongly decayed orbit. This, in turn, requires a massive enough system such that dynamical friction was able to drag it to the inner regions of the Galaxy. This system must have also been rather dense in order to survive the tidal field of the Milky Way and continue to loose orbital energy due to dynamical friction down to an orbit with an apocenter of the order of the Solar circle radius. The current mass of $\\omega$ Cen (5 $10^{6}$ M$_{\\odot}$; Meylan {\\it et al.} 1995) can not generate sufficient dynamical friction to modify its orbit to its current small size (DGvA). Following these arguments, the debris from the massive putative parent galaxy of $\\omega$ Cen may be expected to imprint a kinematical feature in large samples of local, metal-poor stars. The purpose of this investigation is to search for such a feature in the kinematically hot halo. We have used three data sets: the largest, kinematically unbiased sample of metal poor stars ($\\sim 1200$) provided by Beers {\\it et al.} (2000) (hereafter B2000), the sample of globular clusters with measured absolute proper motions (DGvA updated with new distances from Harris 1996, and with a few more clusters; see Dinescu {\\it et al.} 2001), and a small sample of stars with complete kinematics and abundance measurements for O, Na, Mg, Si, Ca, Ti, Cr, Fe, Ni, Y and Ba (Nissen \\& Schuster 1997, hereafter NS97). ", + "conclusions": "We have shown that a distinct population of stars with a metallicity range that inludes the mean metallicity of $\\omega$ Cen, and with $\\omega$ Cen-like phase-space characteristics emerges from the B2000 data. Choosing a metallicity and orbital-parameter range (Section 4) such that we maximize the ``signal'' of this population with respect to the ``noise'' of the halo, we obtain an excess population at 2.3-$\\sigma$ level. By considering the RR Lyrae stars in B2000, we also see this population. We find that the excess RR Lyrae population is predominantly of RRab type, with periods of 0.5 days, and [Fe/H] $\\sim -1.5$. The candidates to have been torn from the system that once contained/was $\\omega$ Cen have one main orbit property: they have a larger eccentricity ($e \\sim 0.8$) (i. e. orbital energy) than that of $\\omega$ Cen ($e = 0.67$). Using the disruption model developed by Johnston (1998) for the orbit of $\\omega$ Cen, we find that trailing tidal debris with orbit characteristics of those of the candidates are found in the Solar neighborhood. We also find that HD 106038, a single, main sequence star with a chemical abundance pattern very similar to that in $\\omega$ Cen stars, in particular enhanced s-process elements (NS97), has $\\omega$ Cen-like orbital properties. Similarly, V716 Oph in B2000 (a BL Her-type variable found in globular clusters of which $\\omega$ Cen is most abundant) has $\\omega$ Cen-like orbital properties, and a metallicity close to the mean metallicity of $\\omega$ Cen. We identify two globular clusters as candidates for belonging to the system that once contained/was $\\omega$ Cen, NGC 362 and NGC 6779. The more metal rich cluster, NGC 362 shows a deficiency in [Cu/Fe] when compared to globular clusters of similar metallicity, a deficiency seen so far only in $\\omega$ Cen stars. {\\bf Acknowledgments}. I am grateful to M\\'{a}rcio Catelan for his suggestions regarding the RR Lyrae stars, and to both M\\'{a}rcio Catelan and Terry Girard for numerous helpful discussions concerning this work. This research has made use of the SIMBAD database, operated at CDS, Strasbourg, France." + }, + "0112/astro-ph0112478_arXiv.txt": { + "abstract": "We present the first calculations to follow the evolution of all stable nuclei and their radioactive progenitors in stellar models computed from the onset of central hydrogen burning through explosion as Type II supernovae. Calculations are performed for Pop I stars of 15, 19, 20, 21, and 25\\,\\Msun using the most recently available experimental and theoretical nuclear data, revised opacity tables, neutrino losses, and weak interaction rates, and taking into account mass loss due to stellar winds. A novel ``adaptive'' reaction network is employed with a variable number of nuclei (adjusted each time step) ranging from $\\sim700$ on the main sequence to $\\gtrsim2200$ during the explosion. The network includes, at any given time, all relevant isotopes from hydrogen through polonium ($Z=84$). Even the limited grid of stellar masses studied suggests that overall good agreement can be achieved with the solar abundances of nuclei between $^{16}$O and $^{90}$Zr. Interesting discrepancies are seen in the 20\\, \\Msun model and, so far, only in that model, that are a consequence of the merging of the oxygen, neon, and carbon shells about a day prior to core collapse. We find that, in some stars, most of the ``$p$-process'' nuclei can be produced in the convective oxygen burning shell moments prior to collapse; in others, they are made only in the explosion. Serious deficiencies still exist in all cases for the $p$-process isotopes of Ru and Mo. ", + "introduction": "\\label{sec:intro} The nucleosynthetic yields of massive stars are important to many areas of astronomical research. Besides the inherent interest in understanding our nuclear origins, the abundances made in supernovae are used to diagnose models for the explosion and as input to still grander models for the formation and chemical evolution of galaxies and the intergalactic medium. They are the target of x-ray observations of supernova remnants and gamma-ray studies of radioactivities in the interstellar medium. Some can be used as cosmochronometers, others power the light curves, still others appear as anomalous abundances found in tiny meteroitic grains in our own solar system. For these reasons, nucleosynthesis calculations have a long history and a sizable community that carries them out. Most recently, nucleosynthesis in massive stars has been studied by \\citet[WW95]{WW95}; \\citet{TNH96,lim00} and others. With this paper, we embark on a new survey, similar to WW95, that will ultimately include stars of many masses and initial metallicities. The characteristics of this new study are improvements in the stellar physics (mass loss rates, opacities, reaction network, etc., \\Sect{comp_proc}) and revisions to nuclear reaction rates (\\Sect{mods}) that have occurred during the last eight years. This first paper particularly addresses recent improvements in nuclear physics. For elements heavier than about silicon, the nuclear level densities are sufficiently high (provided the particle separation energies are not too small) that the statistical - or ``Hauser-Feshbach'' - model can be used. Here, in their maiden voyage, we use rates calculated using the NON-SMOKER code \\citep{rtk97,rt98}. The reaction library, from which the network is drawn, includes all nuclei from the proton-drip line to the neutron-drip line and elements up to and including the actinides \\citep{RATH}. For elements lighter than silicon, where they have been measured, results are taken from the laboratory. Several different compilations are explored. The most critical choices are the rates for $^{12}$C($\\alpha,\\gamma)^{16}$O, $^{22}$Ne($\\alpha$,n)$^{25}$Mg, and $^{22}$Ne($\\alpha,\\gamma)^{26}$Mg. In order to facilitate comparison, we have chosen a constant value equal to 1.2 times that of Buchmann (1996) for the $^{12}$C($\\alpha,\\gamma)^{16}$O rate in {\\sl all} our calculations. For our {\\sl standard} models (defined in $\\S$3.1) we further adopt the lower bound of \\citet{kaepp94} for $^{22}$Ne($\\alpha$,n)$^{25}$Mg \\citep{HWW01}. In future publications we will explore, in greater depth, the consequences of different choices for these rates (for $^{12}$C($\\alpha,\\gamma)^{16}$O, see also \\citealt{WW93,BHW02}). A novel reaction network is employed, unprecedented in size for stellar evolution calculations. The network used by WW95, large in its day, had about 200 nuclides and extended only to germanium. Studies using reaction networks of over 5000 nuclei have been carried out for single zones or regions of stars in order to obtain the $r$-process, e.g., \\citet{CCT85,fre99,kra93}, but ``kilo-nuclide'' studies of nucleosynthesis in complete stellar models (typically of 1000 zones each for 20,000 time steps) have not been done before. We describe in \\Sect{dynet} a dynamically evolving network that adds and subtracts nuclides as appropriate during the star's life to ensure that all significant nuclear flows are contained. Our present survey uses a network that has the accuracy of a fixed network of 2500 isotopes. Section 4 discusses aspects of the stellar evolution that are critical to the nucleosynthesis and \\Sect{results} gives the main results of our survey. We find overall good agreement of our nucleosynthesis calculations with solar abundances for intermediate mass elements (oxygen through zinc) as well as the ``weak component'' of the $s$-process (A $\\ltaprx$ 90), and most of the $p$-process isotopes. However, there is a systematic deficiency of $p$-process isotopes below A $\\approx$ 125 that is particularly acute for Mo and Ru, and around A $\\approx$ 150. Possible explanations are discussed in \\Sect{gamma}. We also find that the nucleosynthesis is at least as sensitive to the stellar model as to the nuclear physics and, in particular, find unusual results for a 20\\,\\Msun model (in the sense that the results differ greatly from both the sun and those at either 19 or 21\\,\\Msun). This is because of the merging of convective oxygen, neon, and carbon shells that occurred well before collapse in that model and not in the others (\\Sect{results}). ", + "conclusions": "Using a nuclear reaction network of unprecedented size, nucleosynthesis has been investigated in several stellar models in the mass range 15 \\Msun to 25 \\Msun. The models include the best currently available nuclear and stellar physics. For the first time, it was also possible to self-consistently follow the $\\gamma$-process up to Bi. Overall good agreement can be achieved with the solar abundances of nuclei between $^{16}$O and $^{90}$Zr. This good agreement is, to first order, independent of the reaction rate set employed; our current standard, \\citet{ang99} or \\citet{HWW01}, though several key nuclear uncertainties are identified. In addition to the well-known need for greater accuracy in the rate for $\\alpha$-capture on \\I{12}C, the rates for \\I{22}{Ne}($\\alpha$,n)\\I{25}{Mg} and $^{22}$Ne($\\alpha,\\gamma)^{26}$Mg are critical. We also urge a re-examination of some of the neutron capture cross sections for the isotopes of nickel. For the $p$-isotopes, two regions of atomic mass are found where those isotopes are underproduced, $92 \\leq A \\leq124$ and $150\\leq A\\leq 165$. It remains unclear whether this deficiency is due to nuclear cross sections, stellar physics, or if alternative (additional) $p$-process scenarios have to be invoked. However, we find that part of the $p$-nuclides may be produced in convective oxygen shell burning during the last hour of the star's life. The remainder is made explosively. Interesting and unusual nucleosynthetic results are found for one particular 20 \\Msun model due to its special stellar structure. This effect, a merging of heavy element shells late in the stars evolution, seems to be confined to a narrow range of masses. In particular it is not seen in 19 and 21 \\Msun models. However, we have explored a very limited set of masses and those only in one spatial dimension (for caveats see Bazan \\& Arnett 1994)." + }, + "0112/astro-ph0112152_arXiv.txt": { + "abstract": "We present optical observations of the fields of two X--ray sources located near the center of the shell-like supernova remnant G266.1--1.2. No objects brighter than R$\\sim$22.5 and B$\\sim$23 are present within the small \\textit{Chandra} error region of \\axj, besides a R$\\sim$17 star that has already been excluded as a possible counterpart. A bright diffuse H$_{\\alpha}$ nebula is present close to the position of the candidate neutron star. ", + "introduction": "The supernova remnant G266.1--1.2 has been reported as a possible $\\gamma$-ray source in the 1.156 MeV line of $^{44}$Ti (Iyudin et al. 1998). The short lifetime ($\\sim$90 yrs) of this isotope, and the relatively small angular size of the remnant would imply an age of only $\\sim$680 years and a small distance d$\\sim$200 pc (Aschenbach et al. 1999). Thus G266.1--1.2 could be the remnant of the closest supernova event to have occurred in recent historical times. However, \\textit{ASCA} observations showed that the X--rays from the SNR shell have a non-thermal spectrum and the fits require a high absorption value (Slane et al. 2001), favoring a distance of $\\sim$1-2 kpc that would place G266.1--1.2 well beyond the Vela SNR (see also Mereghetti \\& Pellizzoni 2001). The \\textit{ASCA} data revealed also a central point source, \\axj , surrounded by diffuse X--ray emission, that was interpreted as the neutron star associated to G266.1--1.2. A \\textit{BeppoSAX} observation (Mereghetti 2001) of the central region of G266.1--1.2 showed the presence of a second source about 3$'$ north of that detected by \\textit{ASCA} and with a harder spectrum. The northern source was named SAX~J0852.0--4615. Since the \\textit{BeppoSAX} error circle of \\axj contained two bright early type stars that might have produced the observed X--ray flux, while no optical counterparts brighter than V$\\sim$15 were visible for \\sax, it was unclear which of the two sources was the most likely neutron star candidate. The puzzle has been recently solved by a \\textit{Chandra} observation that provided an arcsecond position for \\axj (Pavlov et al. 2001). The new error box is incompatible with the two early type stars that were previously considered as possible counterparts, thus confirming that \\axj is the most likely neutron star candidate. \\sax was not detected in the 3 ks long \\textit{Chandra} observation reported by Pavlov et al. (2001). This might be due to variability, or to the hardness of this source, that was detected with \\textit{BeppoSAX} only above 5 keV. A deeper observation with \\textit{XMM-Newton} confirmed the existence of \\sax, with a flux about ten times fainter than that of \\axj (Aschenbach, this conference). Here we present optical observations of the fields of these two X--ray sources. ", + "conclusions": "The deep optical limits for the possible counterparts of \\axj confirm that this is most likely the neutron star remnant associated with G266.1--1.2. An interesting H$_{\\alpha}$ nebula has also been discovered in the data presented here. Emission in the H$_{\\alpha}$ has been detected around a few radio pulsars and is thought to originate in the interstellar medium shocked by the relativistic pulsar wind. These nebulae have either a ''cometary'' shape with the axis of symmetry along the direction of the pulsar transverse motion (e.g., PSR B2224+65 (''Guitar Nebula'', Cordes et al. 1993) or PSR B0740--28 (Stappers et al. these proceedings)) or an arc-like shape (e.g., PSR J0437--4715, Bell et al. 1996). The morphology of the diffuse emission shown in Fig.1 and Fig.2 does not present any obvious connection with the location of the candidate neutron star as determined with \\textit{Chandra}. It is more likely that the nebula is related to the B[e] star Wray 16-30, which is located at the southern end of the nebula. However, its peculiar morphology and the location close to the center of G266.1--1.2 make this nebula a potentially interesting target for more detailed investigations." + }, + "0112/astro-ph0112222_arXiv.txt": { + "abstract": "In this paper I investigate what factors -- both observational and physical -- can change the measured slope of the observed 21cm HI power spectrum. The following effects can make the observed turbulence appear two dimensional rather than three dimensional: 1) if the turbulence is contained in a thin filament or slab; 2) if the medium has a high optical depth; and 3) if any method of observation or analysis is used which effectively limits the emission from the medium under study to a thin slab, for example, by analyzing an individual channel map. Straightforward analysis of data can give misleading or incomplete results if these effects are not taken into account. ", + "introduction": "The 21cm HI line has a column density power spectrum whose slopes are consistent with 2--Dimensional turbulence ($\\alpha \\sim -8/3$) on large spatial scales ($> 0.01$~pc) and narrow velocity ranges (Green 1993; Dickey \\& Crovisier 1983; Lazarian \\& Stanimirovic 2001; Dickey et al. 2001). The slope of this power spectrum is closer to a 3--D, Kolmogorov--like spectrum ($\\alpha \\sim -11/3$) for wider velocity ranges but can still be significantly different than the Kolmogorov value. However, the electron density spectrum is consistent with Kolmogorov--like turbulence. This suggests a) that electrons and neutrals have different turbulent characteristics, or b) that there are effects which change the measured power spectrum slopes. Lazarian \\& Pogosyan (2000) have discussed how turbulent velocity may effect the observed power spectra of HI. However, there are other factors which Lazarian \\& Pogosyan did not discuss which may effect the power spectra. Among these are opacity and filamentary structures in the HI. ", + "conclusions": "" + }, + "0112/astro-ph0112291_arXiv.txt": { + "abstract": "We report upper limits on CO J=2--1 and CO J=5--4 emission from the quasar SDSSp J104433.04-012502.2 at $z=5.73$ from observations made with the Berkeley-Illinois-Maryland-Association Array. Previously reported limits on CO J=6--5 emission (Iwata et al. 2001) were obtained at $z=5.80$, which is now thought to be off by 1\\%, and the observations likely missed the relevant redshifts for molecular gas. The new $3\\sigma$ upper limits on the line luminosities are $L^{'}_{CO}(2-1) < 5.1 \\times 10^{10}$ K~km~s$^{-1}$~pc$^2$ and $L^{'}_{CO}(5-4) < 3.0 \\times 10^{10}$ K~km~s$^{-1}$~pc$^2$, assuming 200~km~s$^{-1}$ linewidth. The CO J=5--4 observations place an upper limit on warm, dense molecular gas mass comparable to amounts derived for some other high redshift quasar systems from detections of this line. The limit on CO J=2--1 emission suggests that excitation bias does not affect this conclusion. In addition, no molecular gas rich companion galaxies are found in a $\\sim1.4$~Mpc field surrounding the quasar. ", + "introduction": "The study of the star formation properties and gas content of galaxies at far cosmological distances is one important step toward understanding galaxy formation and evolution. The quasar SDSSp~J104433.04-012502.2 at $z\\approx5.8$ (hereafter SDSS~1044-0125), discovered by Fan et al. (2000) using the Sloan Digital Sky Survey, is among the highest redshift objects known. Recent observations of this quasar with SCUBA on the JCMT at 850~$\\mu$m detect thermal continuum emission (reported by Iwata et al. 2001), which suggests a large reservoir of dust and therefore also molecular gas. The presence of a substantial gas mass gains support from the apparent X-ray weakness of the quasar, which likely results from heavy intrinsic absorption (Brandt et al. 2001, Mathur 2001). Molecular gas has been detected from at least a dozen $z>2$ objects through CO lines at millimeter wavelengths, and these observations provide important clues to the formation history of galaxies and their relationship to supermassive black holes. A recent search for CO J=6--5 emission from SDSS~1044-0125 by Iwata et al. (2001) reported an upper limit on the inferred molecular gas mass comparable to the detections for some high redshift quasars. Unfortunately, the search was centered at $z=5.80$, the initial redshift estimated by Fan et al. (2000). Recent spectroscopic studies of SDSS~1044-0125 give a more accurate value of $z=5.73\\pm0.01$ (Djorgovski et al. 2001, see also Goodrich et al. 2001), about 1\\% off from the initial estimate. Because of the narrow instantaneous bandwidth available to current millimeter interferometers, the revised redshift falls outside window that was searched for CO J=6--5 emission, and the observations likely missed the relevant redshifts for molecular gas in the quasar host. Since SDSS~1044-0125 has several properties in common with high redshift quasars where CO emission has been detected, (enumerated by Iwata et al. 2001), the revision of the optical redshift determination gives impetus to a new search for molecular gas. A potentially important limitation of searching for emission from CO lines with high rotational quantum numbers, like the J=6--5 transition, is that prevailing physical conditions may be insufficiently extreme to excite these lines. The surprising detection of extended emission in the low excitation CO J=2--1 line towards the quasar APM~08279+5255 at $z=3.91$ (Papadopoulos et al. 2001) suggests that low excitation CO lines can reveal molecular mass reservoirs that are one or two orders of magnitude larger than suggested by observations of high excitation CO lines. In this short paper, we present results of searches for CO J=2--1 and J=5--4 emission from SDSS~1044-0125 using the Berkeley-Illinois-Maryland Array (BIMA)\\footnote{The BIMA array is operated by the Berkeley-Illinois-Maryland Association under funding from the National Science Foundation.} (Welch et al. 1996) that provide new limits on the amount of molecular gas associated with this luminous high redshift quasar. The BIMA 1~cm band receiver system, which was developed primarily for observations of the Sunyaev-Zeldovich effect (Carlstrom, Joy \\& Grego 1996, Grego et al. 2000), provides a unique facility to search for highly redshifted low lying CO lines. For SDSS~1044-0125, the CO J=2--1 line is redshifted to 34 GHz, within the accessible tuning range. The standard digital correlator allows for several times larger velocity coverage than generally available at shorter wavelengths, sufficient to span the uncertainty in the quasar redshift determined from optical lines, as well as the typical kinematic offsets of molecular gas from the redshift derived optically. In addition, the small BIMA array antennas provide a large field of view, which enables imaging the quasar environs over Mpc scales at 34 GHz in a single pointing. ", + "conclusions": "Figure~\\ref{fig:spectrum} shows the CO J=2--1 and CO J=5--4 spectra obtained at quasar position. The velocity binnings for the two lines are 200 and 162~km~s$^{-1}$, respectively. The noise in the CO J=2--1 spectrum is not uniform, and empirical $\\pm1\\sigma$ error bars derived from the rms noise measured from the images are shown for each velocity bin. There are some tantalizing hints of signal in adjacent channels close to the expected velocities, but features with similar (low) significance are present elsewhere in the data, and we do not consider any of these features to be reliable line detections. Figure~\\ref{fig:channels} shows a series of maps with 200~km~s$^{-1}$ width that span the full half power field of view ($6\\farcm6$) for the J=2--1 line. Various attempts at smoothing in both space and frequency did not uncover any significant CO emission in either of the observed transitions. Two mechanisms have been suggested for heating large masses of dust and gas in high redshift quasar systems: (1) high energy photons emitted from gases accreted onto a massive black hole and (2) bursts of star formation. If the dust is heated by the activity of a massive black hole, then bright emission may be expected from high excitation CO lines in a compact region close to the exciting source, from large amounts of warm, dense gas involved in fueling and accretion. The CO J=5--4 line, whose upper energy level lies 88~K above the ground state, requires warm gas ($>30~K$) at high densities ($>10^3$~cm$^{-3}$) to be populated significantly by H$_2$ collisions. Consequently, the upper limits on CO emission from the J=5--4 line constrains primarily the amount of molecular gas with these conditions close to the massive black hole or other powerful heating sources. On the other hand, such extreme physical conditions are not necessarily appropriate for starbursts, which are likely to be distributed over larger spatial scales and involve cooler, more diffuse molecular gas. If the dust is heated by primarily by star formation, then emission in CO J=2--1 line may be a more appropriate tracer of molecular gas content, given that the upper energy level lies just 16~K above ground and the excitation requirements are significantly less stringent. Following Solomon, Downes \\& Radford (1992), we calculate upper limits to CO line luminosities with the expression \\begin{equation} L^{'}_{CO} = 3.25\\times10^{7} S_{CO} \\Delta v \\nu_{obs}^{-2} D_L^2 (1+z)^{-3}~~~~ {\\rm K~km~s}^{-1}~{\\rm pc}^2, \\end{equation} where $S_{CO} \\Delta v$ is the limit on the velocity integrated line flux in Jy~km~s$^{-1}$, $\\nu_{obs}$ is the observing frequency in GHz, and $D_L$ is the luminosity distance in Mpc. The choice of cosmological parameters enters in $D_L$, and we adopt $H_0 = 75$ km~s$^{-1}$, $\\Omega = 1$ and $\\Omega_{\\Lambda} = 0$ for consistency with most work in this field. (An alternative cosmology with $H_0 = 75$ km~s$^{-1}$, $\\Omega = 1$ and $\\Omega_{\\Lambda} = 0.7$ results in $D_L$ larger by a factor of 1.54 for this redshift.) The effective linewidth is not known, but it likely falls in the range 150 to 550~km~s$^{-1}$ found for a large sample of ultraluminous galaxies in the local universe (Solomon et al. 1997). For the $3\\sigma$ flux limit obtained in the more sensitive part of the CO J=2--1 spectrum, assuming a linewidth of 200~km~s$^{-1}$, $L^{'}_{CO}(2-1) < 5.1 \\times 10^{10}$ K~km~s$^{-1}$~pc$^2$. For the $3\\sigma$ flux limit obtained for CO J=5--4, again assuming a linewidth of 200~km~s$^{-1}$, $L^{'}_{CO}(5-4) < 3.0 \\times 10^{10}$ K~km~s$^{-1}$~pc$^2$. If the assumed linewidth were two times larger, then these luminosity limits would be $\\sqrt{2}$ times higher. Conversion of these CO luminosity limits to molecular gas mass limits is fraught with uncertainties. But a simple conversion factor from CO luminosity to H$_2$ mass is commonly taken to be $4.5~M_{\\odot}$ (K~km~s$^{-1}$~pc$^2$)$^{-1}$, the value determined for Milky Way molecular clouds (Sanders, Scoville \\& Soifer 1991). There is evidence from comparisons of luminosity based mass estimates with dynamical mass estimates that the conversion factor may be perhaps five times lower in ultraluminous objects (Downes \\& Solomon 1998). Additional corrections of order unity are also needed to account properly for excitation from the elevated cosmic background radiation at high redshift. Adopting the Galactic conversion factor for CO J=2--1 line luminosity gives a limit on the {\\em cold or diffuse} molecular gas mass of $\\sim2.3\\times10^{11}~M_{\\odot}$ in the SDSS~1044-0125 system. Using the same conversion factor for the CO J=5--4 line luminosity gives a limit on the {\\em warm and dense} molecular gas mass of of $\\sim1.3\\times10^{11}~M_{\\odot}$ in the SDSS~1044-0125 system. These mass limits are comparable to the mass indicated from the detection of CO J=5--4 emission from some $z>4$ quasars, including at least two thought not to be amplified by gravitational lensing. In particular, observations of CO J=5--4 emission from BR1202-0725 at $z=4.7$ (Omont al. 1996, Ohta et al. 1996) and BRI1335-0417 at $z=4.4$ (Guilloteau et al. 1997) indicate molecular gas masses in excess of $10^{11}$~M$_{\\odot}$ (adjusted for the cosmology and CO to H$_2$ conversion factor adopted here). There is no clear physical argument to explain why some quasar environments show CO emission at this sensitivity level while others do not (Guilloteau et al. 1999). In any case, the CO J=2--1 and J=5--4 luminosity limits suggest that the environment of SDSS~1044-0125 does not possess an enormous mass reservoir of either low excitation or high excitation molecular gas. The CO J=2--1 limit is comparable to the amount of molecular gas detected toward the lensed quasar APM~08279+5255, where Papadopoulos et al. (2001) found several CO J=2--1 emission features with total luminosity $6.6\\pm3.1 \\times 10^{11}$ K~km~s$^{-1}$~pc$^2$ attributed to (unlensed) molecular gas rich companion galaxies to the quasar host. For the SDSS~1044-0125 observations, such features would have been contained within one synthesized beam (together with any nuclear emission). The luminosity limit suggests that no comparable population of nearby massive companions is present. Moreover, no significant CO J=2--1 emission features are found within the entire field of view that spans $\\sim1.4$~Mpc, which suggests that such massive cold molecular gas concentrations are rare. Observations of SDSS~1044-0125 with better sensitivity are needed to explore whether smaller but still significant concentrations of low excitation molecular gas are present in the environment of this high redshift quasar." + }, + "0112/astro-ph0112258_arXiv.txt": { + "abstract": "The reduced proper motion diagram (RPMD) for a complete sample of 819 faint ($B \\leq 22.5$) stars with high accuracy proper motions ($\\sigma_{\\mu}\\sim1$ mas yr$^{-1}$) in an area of 0.3 deg$^2$ in the North Galactic Pole field SA57 is investigated. Eight stars with very large reduced proper motions are identified as faint white dwarf candidates. On the basis of larger than 6$\\sigma$ measured proper motions and the lack of photometric variability over a twenty year baseline, we discriminate these white dwarf candidates from the several times more numerous QSOs, which can potentially occupy a similar location in the RPMD. For comparison, less than 4$\\sigma$ proper motions and photometric variability are found in all but one of 35 spectroscopically confirmed QSOs in the same field. While spectroscopic confirmation of their status as white dwarfs is a necessary, but difficult, outstanding task, we discuss the implausibility that these stars could be any kind of survey contaminant. High quality proper motions lend confidence in our ability to separate white dwarfs from subdwarfs in the RPMD. If {\\it bona fide} white dwarfs, the eight candidates found here represent a portion of the white dwarf population that hitherto has remained uninvestigated by previous surveys by virtue of the faint magnitudes and low proper motions of the stars. This faint, low velocity sample represents an increase in the white dwarf sky surface density to $B=22.5$ by an order of magnitude over that found in the previously most complete surveys to this depth. However, because the majority of the stars discovered here are at projected distances of more than a disk scaleheight above the Galactic midplane, their existence does not affect significantly the typical estimates of the local white dwarf density. On the other hand, as distant white dwarf candidates with low, typically thin disk-like transverse velocities ($< 40$ km s$^{-1}$), the newly discovered stars suggest a disk white dwarf scaleheight larger than the values of 250-350 pc typically assumed in assessments of the local white dwarf density (and thought to characterize the Galactic old thin disk in stellar population models). Both a $$ and a more complex maximum likelihood analysis of the spatial distribution of our likely thin disk white dwarfs yield scaleheights of 400-600 pc while at the same time give a reasonable match to the {\\it local} white dwarf volume density found in other surveys (although this good match is a result of the dominance of the one relatively nearby white dwarf in the $1/V_{max}$ density calculation). A high scaleheight persists even if the relatively small sample is pruned of any potential thick disk or halo white dwarfs. While our work is not optimized toward the study of halo white dwarfs as potential MACHO objects, our results do have interesting implications for this hypothesis. We can place some direct constraints (albeit weak ones) on the contribution of halo white dwarfs to the dark matter of the Galaxy. Moreover, the elevated scale height that we measure for the thin disk could alter the interpretation of microlensing results to the extent of making white dwarfs untenable as the dominant MACHO contributor. ", + "introduction": "In Paper I of this series (Majewski 1992), proper motions were determined for nearly a thousand stellar objects in Selected Area 57 (SA57) at the North Galactic Pole to photographic $B_J \\sim 22.5$ and $V_F \\sim 21.5$.\\footnote{$B_J$ is the passband produced by the combination of IIIa-J emulsion and GG385 filter, whereas $V_F$ is the combination of IIIaF + GG495.} Photometric parallaxes were determined for a subsample of 250 stars with $0.3 \\leq B-V \\leq 1.1$ and $U \\leq 21.5$ based on photographic ultraviolet excess measurements. Since, in general, no direct measurement of the surface gravity of each star was readily available, a basic premise of the adopted analysis in Paper I was that the survey stars are on the main sequence. However, it is possible to exploit proper motions to discriminate luminosity classes of some stars through use of the reduced proper motion diagram (RPMD; see also Luyten 1922, Jones 1972a,b, Chiu 1980b, Evans 1992, Knox et al. 1999, hereafter K99; Cooke \\& Reid 2000; Oppenheimer et al. 2001, hereafter O01). White dwarfs, on account of their very high reduced proper motions, should be readily identifiable in the RPMD. This technique confers certain advantages over color searches for white dwarfs; e.g., it is possible to identify cool white dwarfs that are not distinguishable from the more numerous late type field stars using colors alone. Deep searches for faint, cool white dwarfs are important for testing white dwarf cooling models into the regime of Debye crystallization, and, by applying cooling theory in conjunction with the white dwarf luminosity function, to set limits on the star formation history and age of Galactic stellar populations. White dwarfs can also be used as tracers of the density laws of old populations, and white dwarfs are proposed as potentially significant contributors to the dark matter component represented by gravitational microlensing events. Numerous studies have attempted to establish the local density and/or luminosity function of Population I white dwarfs, and especially, recently, at the red end of the white dwarf sequence, due to the interest in cool white dwarfs for both age dating the Galaxy and as a primary source for microlensing candidates. Results for the derived local white dwarf density found among the different surveys still range by a factor of two (Fleming et al. 1986; Jahrei{\\ss} 1987; Liebert, Dahn \\& Monet 1988, hereafter LDM; Boyle 1989; Ruiz \\& Takamiya 1995; Oswalt et al. 1996; Festin 1998; K99; Reid et al. 2001; Ruiz \\& Bergeron 2001). The question of completeness lingers when considering the results of these various surveys. So too does a proper understanding of the density laws appropriate to the samples garnered, since the conversion from a survey list to a local density requires an understanding of the {\\it effective} volume surveyed, i.e., the volume foreshortened by the drop-off in density with distance from the Galactic midplane. Typically, white dwarf studies have {\\it adopted} a standard value for an exponential scaleheight of the disk in such calculations, rather than attempted to {\\it solve} for the density law from their white dwarf samples. This understandable reluctance derives from the relatively limited range of distances probed by complete samples (typically one third to one half of the traditional old disk scaleheight), which limits sensitivity to the form of the density law. Survey incompleteness at distances comparable to a disk scaleheight derives both from photometric {\\it and} astrometric limitations, since proper motions provide the most commonly used means by which to identify white dwarfs (especially those redward of the field star main sequence turn off (MSTO)). Table 1 summarizes the major {\\it astrometric} white dwarf surveys to date (not including studies made from archival survey data, such as the Lowell Proper Motion or Luyten Half-Second catalogues, which we represent by the work of LDM). With the exception of the deep, small area study by Chiu (1980b), these surveys are focused on stars with fairly large proper motions ($\\gtrsim40$ mas yr$^{-1}$). Such a limitation progressively excludes white dwarf populations with ever larger ranges of transverse velocity as a function of distance. Figure 1 shows the limiting distances that are imposed on the detection of white dwarfs as a function of various apparent magnitude and proper motion limits. Figure 2 plots the sky density of detected white dwarfs against both photometric and astrometric limits for the surveys listed in Table 1. Table 1 and Figure 2 demonstrate that the detected white dwarf sky density appears to be more directly correlated with proper motion limits than survey depth. This is an important point, one worth considering given the new emphasis on properly accounting for the total white dwarf density in the foreground of lensed sources. The impressive K99 study, as the deepest, large area survey with the best proper motions to date, provides a benchmark for the present discussion. K99 claim to find no evidence of incompleteness in their survey sample, and that this survey represents the most complete large area sample to date is evidenced by their finding the largest sky surface density of white dwarfs for such a survey to date, 2.07 deg$^{-2}$. Through a variety of arguments, K99 suggest that their proper motion limit of 50-60 mas yr$^{-1}$ provides a reasonable compromise between minimizing spurious detections and maintaining completeness. For example, for bright ($R \\sim 14$) white dwarfs, K99 argue that such a limit is more than enough to detect stars ``having a ({\\it conservative}) tangential velocity of 40 km s$^{-1}$, [which] would have a proper motion of $\\sim 80$ mas yr$^{-1}$\" (emphasis added); clearly for fainter, more distant ($> 100$ pc) examples of stars with similar transverse velocities, however their survey quickly becomes incomplete (Figure 1). It is worth noting at this point that the results of the present analysis will focus on the discovery of white dwarf candidates that are primarily {\\it slower} than K99's ``conservative\" estimate. An important means by which K99 attempt to build their case for completeness is by establishing that their sample mean $(V/V_{max})$ statistic is nearly 0.5 (see \\S 4), where here the $V$'s represent effective volumes {\\it under the assumption of a 300 pc scaleheight}. While K99's investigation of variations in the assumed scaleheight show no significant alteration in their derived luminosity function, the authors do not state how varying the scaleheight affects their assumption of completeness. $$ is primarily a test of uniformity that K99 and others have adapted to a test of completeness. Even K99 admit that obtaining $=0.5$ cannot be regarded as proof of completeness, since an incomplete sample can yield a similar result. While this caveat might be considered all the more prescient if it can be shown that 300 pc is not a proper scaleheight to assume in the effective volume calculations in the first place, in the end, it should be noted that incompleteness in typical magnitude- and proper motion-limited samples may not significantly bias the derivation of the local luminosity function according to the Monte Carlo simulations of Wood \\& Oswalt (1998) and Mendez \\& Ruiz (2001). In this article, we explore the question of the disk white dwarf population density distribution using a low limit proper motion selected sample. We employ reduced proper motions to separate degenerate star candidates from subdwarfs and Population I main sequence stars within the deep proper motion sample of Paper I. Of course, these proper motion techniques have been used in many previous studies, some that explore to similar depth and/or much larger areas of sky than the Paper I sample (Table 1); given our small survey area (0.3 deg$^2$), the volume we probe is much smaller than that explored by most previous surveys, even considering our $B_J=22.5$ magnitude limit. Within the last few years, several new surveys have also reached $B\\sim22.5$, and the K99 survey in particular has the potential, {\\it based on photometric considerations}, to probe much larger volumes than our survey. In practice, however, differences in {\\it astrometric quality} puts the present analysis in a unique niche in parameter space compared to all previous studies: No study in the literature has an astrometric precision comparable to that afforded by the 16 year baseline study using deep photographic exposures on fine-grained emulsions with good plate scale discussed in Paper I. The resultant precision ($\\sim1$ mas yr$^{-1}$ at $B_J=21.5$ and $\\sim1.6$ mas yr$^{-1}$ at $B_J=22.5$) allows for (1) an RPMD that is relatively ``clean\" of astrometric error-induced scatter at the white dwarf locus, minimizing contamination problems, and (2) orders of magnitude smaller proper motion limits. Thus, we can find low proper motion white dwarf candidates that would be missed by all previous surveys, look for white dwarfs at larger distances, and ensure a much higher level of completeness than could be claimed before. The transverse velocity limits of a 1 mas yr$^{-1}$ survey to $B_J=22.5$ as a function of distance are shown in Figure 1; our astrometric advantage for probing to much larger distances than other surveys is clear. As we shall show, this advantage allows us to find a white dwarf sky density at $B_J=22.5$ that is likely an order of magnitude larger than that found by K99, which itself had a density that was larger than found by most previous surveys (see Table 1). The primary contribution to our sky density seems to be from an extended distribution of distant white dwarfs with cold kinematics typical of the thin disk. That the majority of the white dwarfs have projected distances larger than 300 pc, and a third of them beyond 900 pc, suggests the need for a revision of the normally assumed 250-350 pc scaleheight for the thin disk white dwarf population. Thus, a principal endeavor in this contribution is to use this admittedly small, but much less kinematically biased, sample of white dwarfs to define better the vertical distribution of the old disk, a task to which our unusually distant sample has particular leverage. ", + "conclusions": "Under the assumption that all proposed candidates are {\\it bona fide} white dwarfs, the results of the present RPMD analysis of the Paper I survey has yielded a sky density of likely white dwarfs higher by an order of magnitude over the previous most complete samples. Our candidates represent the (expected) low velocity component of the disk white dwarf population excluded by previous proper motion searches. Several of our candidate white dwarfs are fainter and redder than the disk and halo MSTO, so that previous photometric (e.g., UV excess) surveys would not easily have found them. Our results suggest substantial astrometric and photometric incompleteness in previous surveys. Although Wood \\& Oswalt (1998) and Mendez \\& Ruiz (2001) have shown that incompleteness does not significantly bias derivation of the luminosity function or local density (and we concur that the incompleteness we describe here does not likely bias derivation of the local density), they do note that deriving star formation histories from biased samples is highly susceptible to error when the proper motion errors are large ($>$ 100 mas yr$^{-1}$). The tenfold gain in completeness in a large area survey with the photometric depth and astrometric precision of ours would improve the resolution with which star formation histories could be delineated from the white dwarf luminosity function. Such a survey, however, would require long time baseline observations at good plate scale. It is possible that many repeat observations over the course of the Sloan Digital Sky Survey could provide this level of precision for $V > 20$. Of course, HST can achieve such precision, but not over a large area, while the planned FAME and GAIA astrometric missions will deliver the proper motions, but not the depth. Our distant white dwarf candidates, while a small sample, provide leverage on the density law well above the Galactic midplane and suggest a higher white dwarf scaleheight than typically assumed, where the ``old disk scaleheight\" of 250-350 pc falls at the very low end of the ``reasonable\" range of scaleheights derived from the entire candidate white dwarf sample. It is noted that the lowest luminosity white dwarf candidates contribute the highest $V'/V'_{max}$ on average, and when they are excluded the derived scaleheight is lowered, though it remains high compared to white dwarf studies at brighter magnitudes. That we should find white dwarfs to have higher scaleheights than non-degenerate stars seems, at first blush, consistent with white dwarf cooling theory: One might expect the proportion of old stars among white dwarfs to be higher than among unevolved late-type stars and we might expect the highest vertical velocity dispersions for the oldest stars due to secular dynamical heating processes that progressively increase vertical velocity dispersions of stars with time. Yet, the notion of a relatively more heated white dwarf population appears to be at odds with the actual kinematics we measure for our sample of candidate white dwarfs. One might ask what such a dynamically cold population is doing at such large $z$.\\footnote{Note that postulating overestimates of the distances to our stars does not fix the problem because moving the stars closer also make them dynamically colder in the derived transverse velocities.} We note that our velocity data do not fall along the Str\\\"omberg asymmetric drift relation (Binney \\& Merrifield 1998) expected for secularly heated disk stars. A larger sample is needed to check this apparent contradiction. We have found that the method of maximum likelihood provides lower scale heights that are more consistent with (although still generally higher than) those found in the existing literature. Maximum likelihood is also far less sensitive to small number fluctuations. We propose use of this more elegant method of analysis to supplement or replace $$ methods in the future, especially when dealing with small samples. Our statements here must be tempered by two shortcomings of our survey. First of all, it is clear that a larger sample of faint white dwarf candidates with low velocities (requiring more precise proper motions) is needed to better constrain the white dwarf scaleheight, and we hope to increase our sample when additional fields with similar plate material are analyzed. In addition, spectroscopic confirmation of the present and any future deep samples of astrometrically identified white dwarfs would provide much stronger confidence in the interpretation of our results. Radial velocities, if obtainable, will be critical to verifying the kinematics of this sample, but require 6-10 meter class telescopes to obtain. It is hoped that these additional data will help resolve outstanding questions on the spatial distribution of disk white dwarfs." + }, + "0112/astro-ph0112544_arXiv.txt": { + "abstract": "We report on spectral and timing observations of the nearest millisecond pulsar J0437--4715 with the \\chan\\/ X-ray Observatory. The pulsar spectrum, detected up to 7~keV, cannot be described by a simple one-component model. We suggest that it consists of two components, a nonthermal power-law spectrum generated in the pulsar magnetosphere, with a photon index $\\gamma\\approx 2$, and a thermal spectrum emitted by heated polar caps, with a temperature decreasing outwards from 2 MK to 0.5 MK. The lack of spectral features in the thermal component suggests that the neutron star surface is covered by a hydrogen (or helium) atmosphere. The timing analysis shows one X-ray pulse per period, with a pulsed fraction of about 40\\% and the peak at the same pulse phase as the radio peak. No synchrotron pulsar-wind nebula is seen in X-rays. ", + "introduction": "Millisecond (recycled) radio pulsars are distinguished from ordinary pulsars by their very short and stable periods, $P\\la 10$~ms, $\\dot{P}\\sim 10^{-21}-10^{-19}$~s~s$^{-1}$. It is generally accepted that they are very old objects, with spin-down ages $\\tau=P/2\\dot{P}\\sim 10^9-10^{10}$~yr and low surface magnetic fields $B\\propto (P\\dot{P})^{1/2}\\sim 10^8-10^{10}$~G (e.g., Taylor, Manchester, \\& Lyne 1993). Similar to ordinary pulsars, a millisecond pulsar can emit nonthermal X-rays from its magnetosphere, with a hard power-law spectrum and sharp pulsations. In addition to this nonthermal radiation, thermal X-rays can be emitted from the neutron star (NS) surface, provided the surface is hot enough. According to the models of NS thermal evolution (albeit rather uncertain at these old ages), recycled pulsars are too cold (surface temperature $T\\la 0.1$~MK --- see, e.~g., Tsuruta 1998) to be detectable in X-rays. However, their polar caps can be heated up to X-ray temperatures by relativistic particles impinging onto the magnetic poles from the acceleration zones in the magnetosphere. The radio pulsar models (e.g., Cheng \\& Ruderman 1980; Arons 1981; Michel 1991; Beskin, Gurevich, \\& Istomin 1993) predict polar cap radii $\\rpc\\sim (2\\pi R^3/Pc)^{1/2}$ (where $R\\approx 10$~km is the NS radius), i.e., $\\rpc\\sim 1$--5~km for millisecond pulsars, although different models predict quite different polar cap temperatures, in the range of 1--10~MK. Detection of the polar cap thermal radiation would allow one to discriminate between various models of radio pulsars, study the properties of NS surface layers, and constrain the NS mass-to-radius ratio (Pavlov \\& Zavlin 1997; Zavlin \\& Pavlov 1998 [ZP98]). However, just as in the case of ordinary pulsars, this radiation is detectable only if it is not buried under stronger nonthermal radiation. The current theoretical models are not elaborate enough to predict in which (if any) of millisecond pulsars the thermal component can be brighter than the nonthermal one (in particular, both the thermal and nonthermal luminosities are expected to increase with spin-down energy loss $\\dot{E}$, perhaps with different rates). Therefore, we have to rely upon the analysis of X-ray observations to distinguish the thermal and nonthermal components. The X-ray observatories \\ros, \\asca, and \\sax\\/ have detected 11 millisecond pulsars (nearly 1/3 of all X-ray-detected rotation-powered pulsars --- see Becker \\& Pavlov 2001 for a recent review). Five of these pulsars are identified in X-rays only by positional coincidence with the radio pulsars and, due to the low number of recorded counts, provide only crude flux estimates. The radiation from 3 pulsars --- B1821--24 (Saito et al.~1997), B1937+21 (Takahashi et al.~2001), and J0218+4232 (Mineo et al.~2000) --- is clearly nonthermal: their power-law spectra, detected with \\asca\\/ and \\sax\\/ up to energies of 5--10~keV, are very hard, with photon indices $\\gamma\\sim 1$, and their pulse profiles show sharp peaks. Interestingly, these 3 pulsars are characterized by particularly large $\\dot{E}$ values, $\\dot{E}=(2-20)\\times 10^{35}$~erg~s$^{-1}$, and their magnetic fields at the light cylinder, $B_{\\rm lc}=B(R/R_{\\rm lc})^3\\sim 10^6$~G, are close to that of the Crab pulsar. The case for the other 3 pulsars --- J0437--4715 (Becker \\& Tr\\\"umper 1993, 1999 [BT93, BT99]; ZP98), J2124--3358 (BT99), and J0030+0451 (Becker et al.~2000) --- is less certain. These pulsars show broad peaks of X-ray pulsations, but it does not necessarily mean that their radiation is thermal because broad peaks can be produced by nonthermal emission at some viewing angles. High-quality spectra have been recorded for the brightest of these pulsars, J0437--4715, but their interpretation has been controversial --- e.g., ZP98 suggest that the radiation detected with \\ros\\/ and \\euv\\/ can be interpreted as thermal radiation from hot polar caps, whereas BT99 argue that the radiation is nonthermal (see \\S2). To resolve this controversy, the pulsar needed to be observed at energies above the soft \\ros\\/ and \\euv\\/ bands ($E\\ga 2$~keV), and with high spatial resolution to avoid contamination of the pulsar emission by a nearby AGN which compromised the \\asca\\/ and \\sax\\/ data. The {\\sl Chandra} X-ray Observatory provides both the superb spatial resolution and high throughput at higher energies, together with timing capability. In this paper we present the results of our observations of \\psr\\ with \\chan. We start from a summary of the previous results on \\psr\\ in \\S2. The spectral and timing analyses of the \\chan\\/ data are presented in \\S3 and \\S4. Implication of the results are discussed in \\S5. ", + "conclusions": "The \\chan\\/ observations of \\psr\\ have allowed us to perform the spatial, timing and spectral analyses of the new data collected with high angular and spectral resolution in an extended energy range. The principal new result of the ACIS observation is the measurement of the pulsar's X-ray spectrum at higher energies, up to 7~keV. Among the formally acceptable fits of the combined \\ros\\/ PSPC and \\chan\\/ ACIS spectra, only the broken PL model corresponds to a purely nonthermal pulsar radiation in the X-ray band. However, the break energy of about 1.1 keV is 5--6 orders of magnitude lower than those observed in other pulsars. Such a difference looks too large to be explained by a lower magnetic field in the radiating region --- it would also require a lower break energy in the spectrum of radiating electrons/positrons, compared to ordinary pulsars. Moreover, extrapolation of the broken PL model to the optical $B$ and $V$ bands (assuming the same slope as in the X-ray range below 1.1~keV) predicts optical magnitudes, $m_B=19.6\\pm 1.4$ and $m_V=19.1\\pm 1.4$ (for extinction coefficients $A_B=0.3$, $A_V=0.2$ --- Danziger, Baade, \\& della Valle 1993), much brighter than those detected from the white dwarf companion, $m_B=22.1\\pm 0.1$ and $m_V=20.9\\pm 0.1$ (Bailyn 1993; Danziger et al.~1993; Bell et al.~1993). Thus, we do not consider the broken PL model to be a plausible interpretation. The alternative model involves {\\sl two components} of different origin --- thermal and nonthermal. The nonthermal component originates in the pulsar magnetosphere\\footnote{Another source of unresolvable nonthermal radiation could be the shocked pulsar wind near the white dwarf companion (e.g., Arons \\& Tavani 1993). However, at the distance of $a_p=1.5\\times 10^{11}$ cm from the pulsar (van Straten et al.\\ 2001), the companion intercepts only a fraction $\\sim 6\\times 10^{-4}$ of the wind (assuming the wind is approximately isotropic), too small to explain the observed nonthermal luminosity.}, whereas the thermal component is emitted from hot polar caps on the NS surface. Depending on assumption about the polar cap temperature distribution, one gets different relative contributions from these two components. In the model with uniformly heated polar caps, the nonthermal component described as a PL of photon index $\\gamma=2.7-2.9$ provides about 80\\% in the X-ray flux and dominates at energies below 0.6 keV and above 2.7~keV. However, this model encounters the same problems in the EUV and optical bands as the broken PL (the steeper slope of this PL component predicts even higher fluxes in the $B$ and $V$ bands). If we assume a more plausible polar cap model with temperature decreasing outwards from the cap center, then the thermal component becomes dominant between 0.06 keV and 2.5 keV, providing some 75\\% of the X-ray flux, while the PL component of $\\gamma=1.6-2.5$ dominates outside this band. In addition to a more realistic temperature distribution, the latter model is well consistent with the \\euv\\/ data and yields estimates on the hydrogen column density in agreement with the indirect measurements. As this PL component is fainter than in the two other models, its extension to the optical falls below the observed radiation of the white dwarf companion for photon indices $\\gamma < 1.9$. This allows us to predict that the PL component should be observable in the UV (particularly, far-UV) range where it is brighter than the Wien tail of the white dwarf spectrum, assuming there is no turnover of the nonthermal spectrum between the UV and soft-X-ray energies. On the other side of the X-ray band, extrapolation of the PL component to the gamma-ray energies above 100~MeV predicts a photon flux $f<2\\times 10^{-8}$~s$^{-1}$~cm$^{-2}$ ($\\gamma>1.6$), below the upper limit, $f < 1.5\\times 10^{-7}$~s$^{-1}$~cm$^{-2}$, obtained from the {\\sl CGRO} EGRET observations (Fierro et al.~1995). We emphasize that these models require the thermal radiation to be emitted from hydrogen (or helium) NS atmosphere. The high spectral resolution of the ACIS data rules out an atmosphere comprised of heavier chemical elements. The HRC-S observation of \\psr\\ has demonstrated the \\chan\\/ timing capability at a millisecond level. The HRC-S pulse profile looks narrower, and the pulsed fraction is somewhat higher, than those obtained in the earlier \\ros\\/ and \\euv\\/ observations at lower energies, which could be explained by the properties of the thermal radiation from polar caps covered with a hydrogen or helium atmosphere. On the other hand, the shape of the profile is clearly asymmetric, with a longer rise and faster decay, which cannot be explained by a simple axisymmetric temperature distribution. Relativistic effects (particularly, the Doppler boost) should lead to a different asymmetry --- a faster rise and longer trail (Braje \\& Romani 2000; Ford 2000). The analysis of HRC-S data demonstrates, for the first time, that the phase of the X-ray pulse virtually coincides with that of the radio pulse. If, as we suggest, the main contribution to the HRC-S band is due to the thermal polar cap radiation, and if the pulsar radio beam is directed along the magnetic axis, then the radio emission must be generated close to the NS surface --- e.g., the time difference of $<0.1$ ms between the X-ray and radio phases corresponds to a distance of $<30$ km, much smaller than the light cylinder radius, $R_{\\rm lc}=275$ km. Alternatively, if the radio emission is generated at a higher altitude, the combination of field-line sweepback and aberration must contrive to cancel the radial travel-time difference. The \\chan\\/ observations show no sign of an X-ray PWN that could accompany the bow-shock revealed by the H$_\\alpha$ observations. Three-sigma upper limits on the PWN brightness (in counts~arcsec$^{-2}$) can be estimated as $3(b/A)^{1/2}$, where $b$ is the background surface brightness ($b=0.51$ and 0.28 counts~arcsec$^{-2}$ for the ACIS-S and HRC-I images, respectively), and $A$ is the PWN area (we will scale it as $A=1000 f_A$ arcsec$^2$, assuming that a typical transverse size of the PWN somewhat exceeds the stand-off distance, $10''$, of the bow shock). For a power-law PWN spectrum with a photon index $\\gamma=1.5$--2 (similar to those observed from other PWNe), these upper limits correspond to the PWN intensities $I_{x,{\\rm pwn}}<(1.3$--$1.8)\\times 10^{-17} f_A^{-1/2}$ and $I_{x,{\\rm pwn}}<(3.6$--$5.7)\\times 10^{-17} f_A^{-1/2}$ erg cm$^{-2}$ s$^{-1}$ arcsec$^{-2}$, for the ACIS-S and HRC-I, respectively, in the 0.1--10 keV range. The corresponding upper limits on the PWN X-ray luminosity, $L_{x,{\\rm pwn}}\\approx 4\\pi d^2 A I_{x,{\\rm pwn}}$ are much smaller than the rotational energy loss rate, $\\dot{E}=3.8\\times 10^{33}$ erg s$^{-1}$ --- e.g., $L_{x,{\\rm pwn}} < (3.0$--$4.2)\\times 10^{28} f_A^{1/2}$ erg s$^{-1}$ for the more sensitive ACIS-S limit. The low upper limits on the PWN luminosity in X-rays can be simply explained by a low magnetic field in the PWN region, expected for a particle-dominated pulsar wind. The shock in the relativistic pulsar wind should be located just interior to the observed H$_\\alpha$ bow shock (Arons \\& Tavani 1993). When the wind electrons pass through the shock, their directions of motion become ``randomized'', and their synchrotron radiation may result in an X-ray nebula, provided the electron energies and the magnetic field are high enough in the post-shock region. The pre-shock magnetic field can be estimated as $B_1=[\\dot{E}/(f_\\Omega r_s^2c)]^{1/2} [\\sigma/(1+\\sigma)]^{1/2}=18\\,f_\\Omega^{-1/2} [\\sigma/(1+\\sigma)]^{1/2}~\\mu$G, where $r_s\\approx 2\\times 10^{16}$ cm is the stand-off distance corresponding to $10''$ at $d=140$ pc, $f_\\Omega=\\Delta\\Omega/(4\\pi)\\leq 1$ is the collimation factor of the wind, and the ``magnetization parameter'' $\\sigma$ is the ratio of the Poynting flux to the kinetic energy flux. The maximum value of the post-shock magnetic field, $B\\simeq B_1\\simeq 18\\, f_\\Omega^{-1/2}~\\mu$G, is obtained for $\\sigma\\gg 1$. However, according to Kennel \\& Coronity (1984; KC84 hereafter), a significant fraction of the total energy flux upstream can be converted into (observable) synchrotron luminosity downstream only if $\\sigma \\lapr 0.1$ (e.g., these authors estimate $\\sigma\\approx 0.003$ for the Crab pulsar). For $\\sigma\\ll 1$, the post-shock magnetic field is $B\\simeq 3(1-4\\sigma) B_1\\simeq 53\\,f_\\Omega^{-1/2} \\sigma^{1/2}(1-4.5\\sigma)~\\mu$G (e.g., from eqs.\\ [4.8] and [4.11] of KC84, $B$ increases from 3\\,$\\mu$G to 12\\,$\\mu$G when $\\sigma$ increases from 0.003 to 0.1, at $f_\\Omega=1$). Such low magnetic fields in the shocked wind strongly limit the maximum energy, $m_ec^2 \\Gamma_{\\rm max}$, of radiating electrons and, consequently, the maximum frequency $\\nu_{\\rm max}$ of the synchrotron radiation. Since the Larmor radius of most energetic electrons, $r_{L}=1.7\\times 10^8\\, \\Gamma_{\\rm max} B_{-5}^{~~-1}~{\\rm cm}$, cannot exceed the shock radius $r_s$ substantially, we obtain $\\Gamma_{\\rm max} < 10^8\\, f_s B_{-5}$, $h\\nu_{\\rm max} \\sim (heB/4\\pi m_e c)\\Gamma_{\\rm max}^2 < 0.6\\, f_s^2 B_{-5}^{~~3}$ keV, where $B_{-5}=B/(10\\, \\mu{\\rm G}$), $f_s\\equiv r_L/r_s\\sim 1$. If, for instance, $B<5 f_s^{-2/3}~\\mu$G (i.e., $\\sigma < 0.01\\, f_s^{-4/3}$ in the KC84 model), the synchrotron emission at the bow shock is not expected to be seen in X-rays. Despite the superior quality of the \\chan\\/ data, which have allowed us to detect the hard tail of the pulsar's spectrum and pinpoint the absolute phase of the X-ray pulse, there still remain some open problems. Although our analysis strongly favors the thermal+nonthermal interpretation, it is still unclear which of the two components dominates in the X-ray radiation of \\psr. To establish the relative contributions of these components, energy-resolved timing and time-resolved spectral analysis are needed, which, hopefully, will be possible with the forthcoming {\\sl XMM-Newton} data." + }, + "0112/astro-ph0112402_arXiv.txt": { + "abstract": "We consider the metallicities and kinematics of nearby stars known to have planetary-mass companions in the general context of the overall properties of the local Galactic Disk. We have used Str\\\"omgren photometry to determine abundances for both the extrasolar-planet host stars and for a volume-limited sample of 486 F, G and K stars selected from the Hipparcos catalogue. The latter data show that the Sun lies near the modal abundance of the disk, with over 45\\% of local stars having super-solar metallicities. Twenty of the latter stars (4.1\\%) are known to have planetary-mass companions. Using that ratio to scale data for the complete sample of planetary host stars, we find that the fraction of stars with extrasolar planets rises sharply with increasing abundance, confirming previous results. However, the frequency remains at the 3-4\\% level for stars within $\\pm0.15$ dex of solar abundance, and falls to $\\sim1\\%$ only for stars with abundances less than half solar. Given the present observational constraints, both in velocity precision and in the available time baseline, these numbers represent a lower limit to the frequency of extrasolar planetary systems. A comparison between the kinematics of the planetary host stars and a representative sample of disk stars suggests that the former have an average age which is $\\sim60\\%$ of the latter. ", + "introduction": " ", + "conclusions": "" + }, + "0112/astro-ph0112128_arXiv.txt": { + "abstract": "This paper is the second in a series devoted to examining the multi-wavelength properties of supernova remnants (SNRs) located in nearby galaxies. We consider here the resident SNRs in the nearby Sculptor Group Sd galaxy NGC 7793. Using our own Very Large Array (VLA) radio observations at 6 and 20~cm, as well as archived $\\it{ROSAT}$ X-ray data, previously published optical results and our own H$\\alpha$ image, we have searched for X-ray and radio counterparts to previously-known optically-identified SNRs, and for new previously unidentified SNRs at these two wavelength regimes. Consistent with our prior results for NGC 300, only a tiny minority of the optically-identified SNRs have been found at another wavelength. The most noteworthy source in our study is N7793-S26, which is the only SNR in this galaxy that is detected at all three wavelengths (X-ray, optical and radio). It features a long ($\\sim$ 450 parsecs) filamentary morphology that is clearly seen in both the optical and the radio images. N7793-S26's radio luminosity exceeds that of the Galactic SNR Cas A, and based on equipartition calculations we determine that an energy of at least 10$^{52}$ ergs is required to maintain this source. Such a result argues for the source being created by multiple supernova explosions rather than by a single supernova event. A second optically-identified SNR, N7793-S11, has detectable radio emission but is not detected in the X-ray. A radio-selected sample of candidate SNRs has also been prepared by searching for coincidences between non-thermal radio sources and regions of H$\\alpha$ emission in this galaxy, and this search has yielded five new candidate radio SNRs, to be added to the 28 SNRs that have already been detected by optical methods. A complementary search for new candidate X-ray SNRs has also been conducted by searching for soft-spectrum sources ($\\it{kT}$ $<$ 1 keV) that are coincident with regions of H$\\alpha$ emission. That search has yielded a candidate X-ray SNR which is coincident with one (and possibly two) of the candidate radio SNRs, but considerable diffuse X-ray emission throughout the disk of NGC 7793 reduces the efficacy of the search. Like NGC 300, very little overlap in identifications is seen between the SNRs found through X-ray, optical and radio methods, and such a result argues for the role played by distance-dependent selection effects in determining the detectability of SNRs. In addition, we find that the density of the ambient interstellar medium (ISM) surrounding SNRs significantly impacts the spectral characteristics of the SNRs in this galaxy, consistent with surveys of the SNR populations in other galaxies. ", + "introduction": "This paper is the second in a series devoted to the multi-wavelength study of supernova remnants (SNRs) in nearby galaxies. In our previous paper \\citep*[hereafter referred to as Paper I]{Pannuti00} we analyzed observations made at the X-ray, optical and radio wavelengths of the nearby Sculptor Group Sd galaxy NGC 300. We sought to determine the X-ray and radio properties of the 28 SNRs identified previously in that galaxy through optical search techniques, namely H$\\alpha$ and [\\ion{S}{2}] narrow-band imaging \\citep*[hereafter referred to as BL97]{DDB80, BL97}, and in addition, we searched for new candidate SNRs at the X-ray and radio wavelengths to complement this prior optical work. Our search yielded sixteen new candidate X-ray and radio SNRs, later reduced to fifteen by the recent work of \\citet{RP01}, and the total number of SNRs and candidate SNRs in NGC 300 is now 43. We found very little overlap between the three sets of selected candidates, and we interpret this to indicate that a multiple-wavelength approach is necessary to detect a maximum number of candidate SNRs in a particular galaxy. We also hypothesized that the limited overlap between the selected sets of candidate SNRs indicated selection effects inherent in each type of survey: optical surveys are biased toward the detection of SNRs located in regions of low density and corresponding low optical confusion, while X-ray and radio surveys have the opposite bias and favor the detection of SNRs located in regions of high density. \\par In this paper, we examine another Sculptor Group Sd galaxy, NGC 7793, and once again we consider observations made at the X-ray, optical and radio wavelengths. Following the paradigm set by our previous work (\\citeauthor{Pannuti00}), our intent is to determine the X-ray and radio properties of the previously-known optically-identified SNRs as well as search for new candidate SNRs at those two wavelengths. Salient properties of NGC 7793 are listed in Table \\ref{7793Props}: \\citet{PC88} measured a distance to this galaxy of 3.38 Mpc and classified it as a member of the Sculptor Group of galaxies. Because of its proximity and its low inclination angle of 50$^{\\circ}$ \\citep{T88}, this galaxy makes an excellent choice for the study of galactic properties. These studies include an optical survey for resident SNRs (\\citeauthor{BL97}), analyses of its HI content \\citep{CP90}, its surface photometry \\citep{C85} and its radio continuum properties \\citep{H86}. Its X-ray properties have also been the subject of X-ray analysis, based on observations performed with the $\\it{Einstein}$ satellite \\citep{F92} and later observations performed with the $\\it{ROSAT}$ satellite \\citep*[hereafter referred to as RP99]{RP99}. In Section \\ref{NGC7793ObsSection}, we describe the observations of this galaxy and data reduction at each wavelength, beginning with the radio (Section \\ref{NGC7793RadioObsSubSection}), followed by the optical (Section \\ref{NGC7793OptObsSubSection}) and concluding with the X-ray (Section \\ref{NGC7793XrayObsSubSection}). We discuss the multi-wavelength properties of the optically-selected SNRs in Section \\ref{NGC7793OptSNRsSection}, and then present the new candidate SNRs selected at the radio and X-ray wavelengths in Sections \\ref{NGC7793RadioSNRsSection} and \\ref{NGC7793XraySNRsSection}. A discussion of our findings in this work is presented in Section \\ref{DiscussionSection} and finally our conclusions are given in Section \\ref{ConclusionsSection}. ", + "conclusions": "} We have presented a multi-wavelength search and analysis of the SNR population in NGC 7793. The results and conclusions of this work can be summarized as follows: \\par 1) NGC 7793, a nearby spiral galaxy, has been observed in the X-ray, optical and radio wavelengths to analyze its resident SNR population. This analysis has examined X-ray observations of the galaxy made by the ${\\it ROSAT}$ ${\\it PSPC}$ instrument, an H$\\alpha$ image of the galaxy and new 6 and 20~cm observations made with the VLA. We have analyzed both the X-ray and radio spectral properties of the SNRs previously identified in the optical by \\citeauthor{BL97}, and in addition we have searched for new candidate X-ray and radio SNRs. \\par 2) N7793-S26 is the only optically-identified SNR that possesses both X-ray and non-thermal radio emission. The extreme radio luminosity of this object and its large filamentary structure have led to speculation about its creation. A study of the energetics of this SNR suggest that it was created by multiple supernova explosions rather than a single supernova event. One other optically-identified SNR, N7793-S11, possesses non-thermal radio emission but no X-ray emission. Consistent with prior studies of the SNR populations in the galaxies NGC 300 and NGC 6946, X-ray and radio emission from optically-identified SNRs are in general not detected. This lends more support to the hypothesis that searches conducted for SNRs at each wavelength -- X-ray, optical and radio -- all possess inherent biases. We also find additional evidence that the density of the ambient ISM surrounding an SNR plays a critical role in dictating the SNR's spectral characteristics. \\par 3) A search for non-thermal radio sources at 6~cm and 20~cm (with a minimum detection level of 3$\\sigma$) that are close to or within HII regions has yielded five candidate radio SNRs. Of these five sources, one (NGC 7793-R3) possesses an X-ray counterpart at the 3$\\sigma$ level, namely P10, from the listing of X-ray sources in this galaxy that was prepared by \\citeauthor{RP99}. Another candidate radio SNR, NGC 7793-R4, may also contribute X-ray flux to P10. \\par 4) A search for candidate X-ray SNRs has not revealed any new sources in addition to the SNRs already found through radio and optical surveys. The search for SNRs at this wavelength is complicated by the presence of considerable diffuse X-ray emission throughout the entire disk of this galaxy. \\par 5) The multi-wavelength campaign has added five new candidate SNRs to the 28 previously identified through optical methods. While the number of new detected SNRs is noticeably lower than the number of new SNRs found in our study of NGC 300 as presented in \\citeauthor{Pannuti00}, we feel that this result is linked to the larger distance to NGC 7793 than NGC 300. Because the X-ray and radio observations of NGC 7793 did not have improved sensitivities to compensate for the increased distance to this galaxy compared to NGC 300, both surveys can only sample more luminous portions of the X-ray and radio SNR population in NGC 7793 compared to NGC 300." + }, + "0112/hep-ph0112247_arXiv.txt": { + "abstract": "\\PRE{\\vspace*{.1in}} If extra spacetime dimensions and low-scale gravity exist, black holes will be produced in observable collisions of elementary particles. For the next several years, ultra-high energy cosmic rays provide the most promising window on this phenomenon. In particular, cosmic neutrinos can produce black holes deep in the Earth's atmosphere, leading to quasi-horizontal giant air showers. We determine the sensitivity of cosmic ray detectors to black hole production and compare the results to other probes of extra dimensions. With $n \\ge 4$ extra dimensions, current bounds on deeply penetrating showers from AGASA already provide the most stringent bound on low-scale gravity, requiring a fundamental Planck scale $M_D > 1.3 - 1.8~\\tev$. The Auger Observatory will probe $M_D$ as large as 4 TeV and may observe on the order of a hundred black holes in 5 years. We also consider the implications of angular momentum and possible exponentially suppressed parton cross sections; including these effects, large black hole rates are still possible. Finally, we demonstrate that even if only a few black hole events are observed, a standard model interpretation may be excluded by comparison with Earth-skimming neutrino rates. ", + "introduction": "Tiny black holes (BHs) can be produced in particle collisions with center-of-mass energies above the fundamental scale of gravity~\\cite{Amati:1987wq,'tHooft:rb}, where they should be well-described semi-classically and thermodynamically~\\cite{Hawking:1975sw}. In conventional 4-dimensional theories, {\\em viz.}, where the Planck scale $\\sim 10^{19}~\\gev$ is fundamental and the weak scale $\\sim 1$ TeV is derived from it via some dynamical mechanism, the study of such BHs is far beyond the realm of experimental particle physics. Over the last few years, however, physicists have begun exploring an alternative approach to the longstanding gauge hierarchy problem, wherein the weak scale becomes the fundamental scale of nature and the Planck scale is derived from this, with the hierarchy in scales a consequence of large or warped extra dimensions~\\cite{Antoniadis:1990ew,Randall:1999ee}. If this is the case, the fundamental scale of gravity can be ${\\cal O}$(TeV), and BH production and evaporation may be observed in collisions of elementary particles~\\cite{Banks:1999gd,Emparan:2000rs,% Giddings:2000ay,Giddings:2001bu,Dimopoulos:2001hw}. If gravity indeed becomes strong at the TeV scale, ultra-high energy cosmic rays provide a powerful opportunity to probe BH production at super-Planckian energies~\\cite{Feng:2001ib}. Cosmic rays with energies $\\sim 10^{19}~\\ev$ have been observed~\\cite{Nagano:ve}. They interact in the Earth's atmosphere and crust with center-of-mass energies $\\sim 100~\\tev$, far beyond the reach of present and planned man-made colliders. These cosmic rays may therefore produce BHs, allowing cosmic ray detectors to test the existence of TeV-scale gravity and extra dimensions by searching for evidence of BH production~\\cite{Feng:2001ib,Anchordoqui:2001ei,% Emparan:2001kf,Ringwald:2001vk,Uehara:2001yk}. A particularly promising signal is provided by ultra-high energy cosmic neutrinos, which may produce BHs with cross sections two or more orders of magnitude above their standard model (SM) interactions. These BHs will decay promptly in a thermal distribution of SM particles. Of the order of a hundred BH events may be detected at the Auger Observatory~\\cite{Feng:2001ib} as quasi-horizontal, deeply penetrating showers with distinctive properties~\\cite{Anchordoqui:2001ei}. The possibility of BH production by cosmic rays supplements possible sub-Planckian signatures of low-scale gravity~\\cite{Nussinov:1999jt,Jain:2000pu,Tyler:2001gt,% Alvarez-Muniz:2001mk,Sigl:2001th}. In this article we extend previous work to derive bounds from the non-observation of BH-initiated showers in current data at the Akeno Giant Air Shower Array (AGASA). We also extend previous analyses of BH discovery prospects at Auger, and discuss in detail the possibility of distinguishing BH events from SM events. A preliminary version of some of these results was presented in Ref.~\\cite{GAP}. We begin in Sec.~\\ref{sec:limits} with an overview of TeV-scale gravity. We collect and review existing bounds on the fundamental Planck scale in a uniform convention. In Sec.~\\ref{sec:BH} we discuss semiclassical BH production, including the effects of angular momentum and the production of Kerr BHs, as well as the proposed exponential suppression advocated by Voloshin~\\cite{Voloshin:2001vs,Voloshin:2001fe}. This is followed in Secs.~\\ref{sec:flux} and \\ref{sec:acceptance} by detailed discussions of cosmogenic neutrino fluxes and ground array experiments, respectively. Our results for event rates and new limits on the scale of higher-dimensional gravity are presented in Secs.~\\ref{sec:AGASA} and \\ref{sec:Auger}. We begin with current data from AGASA. The AGASA Collaboration has already reported no significant signal for neutrino air showers during an observation time (live) of 1710.5 days~\\cite{agasa}. Given the standard assumption of a geometric black hole cross section, we find that this data implies the most stringent bound on the fundamental Planck scale to date for $n \\ge 4$ extra dimensions, exceeding limits derived~\\cite{Ringwald:2001vk} from Fly's Eye data~\\cite{Baltrusaitis:mt} and also more stringent than the constraints from graviton emission and exchange obtained by the LEP~\\cite{Pagliarone:2001ff} and D\\O~\\cite{Abbott:2000zb} Collaborations. In Sec.~\\ref{sec:Auger} we then consider the prospects for BH production at the Auger Observatory. Tens of black holes may be observed per year; conversely, non-observation of BHs will imply bounds as large as 4 TeV on the fundamental Planck scale. In Sec.~\\ref{sec:skimming} we note that comparison to Earth-skimming neutrino event rates~\\cite{Bertou:2001vm,Feng:2001ue,Kusenko:2001gj,Domokos:1997ve} allows one to distinguish BH events from SM events. This point was noted already in Ref.~\\cite{Feng:2001ib}, but was not considered in Ref.~\\cite{Ringwald:2001vk}, leading to weaker conclusions. Here, we consider this point quantitatively and find that, even with a handful of BH events, a SM explanation may be excluded based on event rates alone. If seen, black holes created by cosmic rays will provide the first evidence for extra dimensions and TeV-scale gravity, initiating an era of detailed study of black hole properties at both cosmic ray detectors and future colliders, such as the LHC~\\cite{Giddings:2001bu,Dimopoulos:2001hw,Cheung:2001ue,% Park:2001xc,Rizzo:2001dk,Dimopoulos:2001qe,Hossenfelder:2001dn}. Our conclusions are collected in Sec.~\\ref{sec:conclusions}. ", + "conclusions": "\\label{sec:conclusions} In this work we have shown that cosmic ray observations in the recent past (AGASA) and in the near future (Auger) provide extremely sensitive probes of low-scale gravity and extra dimensions. We have focused on the production of TeV-scale BHs resulting from collisions of ultra-high energy cosmic neutrinos in the Earth's atmosphere, and have considered the impact of various theoretical issues in the determination of the BH production cross section. In particular, mass shedding, the production of BHs with non-zero angular momentum, and a possible enhancement of the BH cross section can be expected to give minor perturbations. The exponential suppression proposed by Voloshin is more significant, but large and observable BH event rates are still possible. More specifically, in the case of $n$ extra spatial dimensions compactified on an $n$-torus with a common radius, we have found the following: \\begin{itemize} \\item{Present bounds on atmospheric BH production imply 95\\% CL lower limits on the fundamental Planck mass of $\\md\\ge 1.3-1.5~\\tev$ for $n=4$, rising to $\\md\\ge 1.6-1.8~\\tev$ for $n=7$. These bounds follow from the non-observation of a significant excess of deep, quasi-horizontal showers in 1710.5 days of running recently reported by the AGASA Collaboration~\\cite{agasa}. The absence of a deeply-penetrating signal in the Fly's Eye data~\\cite{Baltrusaitis:mt} also implies lower bounds on $M_D$. These are consistently weaker, however. For example, for $n=6$, $\\xmin=1$, and the same (PJ) flux we have used, Ringwald and Tu find $M_D > 900~\\gev$~\\cite{Ringwald:2001vk}. We find this difference to be significant: the AGASA and Fly's Eye constraints rely on identical theoretical assumptions, and given the scaling in \\eqref{scaling}, a factor of 2 difference in $\\md$ bounds corresponds to a factor of more than 4 in acceptance or, equivalently, running time. The AGASA limits derived here exceed the D\\O\\ bound $\\md\\agt 0.6-1.2$ TeV, where the variation reflects uncertainty from the choice of ultraviolet cutoff for graviton momenta transverse to the brane. The cosmic ray limits are subject to a separate set of uncertainties, discussed at length above, but follow from conservative evaluations of the neutrino flux and experimental aperture, and $\\xmin=1$. For $\\xmin=3$, these limits are somewhat reduced, but still generally exceed the Tevatron bounds. The cosmic ray bounds from AGASA therefore represent the best existing limits on the scale of TeV-gravity for $n\\ge 4$ extra spatial dimensions. A summary of the most stringent present bounds on $\\md$ for $n \\ge 2$ extra dimensions is given in Fig.~\\ref{fig:summary}. } \\begin{figure}[tbp] \\postscript{summary.eps}{0.56} \\caption{Bounds on the fundamental Planck scale $\\md$ from tests of Newton's law on sub-millimeter scales, bounds on supernova cooling and neutron star heating, dielectron and diphoton production at the Tevatron, and non-observation of BH production at AGASA. Future limits from the Auger ground array, assuming 5 years of data and no excess above the SM neutrino background, are also shown. The range in Tevatron bounds corresponds to the range of brane softening parameter $\\Lambda/\\md =0.5-1$. The range in cosmic ray bounds is for $\\xmin=1-3$. See text for discussion. } \\label{fig:summary} \\end{figure} \\item{The reach of AGASA will be extended significantly by the Auger Observatory. If no quasi-horizontal extended air shower events are observed in 5 years (beyond the expected two SM neutrino events supplemented by as many as 10 hadronic background events), Auger will set a limit of $\\md\\agt 3~\\tev$, at 95\\% CL, for $n \\ge 4$. Even in the case where the cross section is decreased by the exponential suppression factor in \\eqref{sigmasupp}, a bound $\\md\\agt$ 2 TeV may be found under the same background assumptions. } \\item{Conversely, given the large reach of Auger, tens of BH events may be observed per year. We have discussed in some detail how combined measurements of quasi-horizontal air showers and Earth-skimming $\\nu_{\\tau}\\rightarrow \\tau$ events may be used to identify new neutrino interactions beyond the SM, even with complete uncertainty about the incident neutrino flux. In the case of BH production, the quasi-horizontal event rate is enhanced, while the Earth-skimming rate is suppressed, since BH production in the Earth acts as an absorptive channel, depleting the SM rate. With counting experiments alone, one can therefore exclude a SM interpretation of BH events, and may distinguish BH events from almost all other possible forms of new physics.} \\end{itemize} In conclusion, in the next several years prior to the analysis of data from the LHC, super-Planckian BH production from cosmic rays provides a promising probe of extra dimensions. Searches for BH-initiated quasi-horizontal showers in the Earth's atmosphere at AGASA provide the most stringent bounds on low-scale gravity at present, and the Auger Observatory will extend this sensitivity to fundamental Planck scales well above the TeV scale." + }, + "0112/astro-ph0112490_arXiv.txt": { + "abstract": " ", + "introduction": " ", + "conclusions": "I conclude that it is practicable to measure $p\\overline{p}$ annihilation spectra separately from underlying power-law backgrounds if the appropriate selection criteria are applied. By hypothesis, any source emitting this spectral feature must be artificial; it is extremely difficult to imagine any other possibility except new laws of physics (\\S 1). When applied to the EGRET data the method would have detected a steady $p\\overline{p}$ annihilation spectrum down to levels $\\sim 2 \\times 10^{-8}$ photon/(cm$^{2}$ s), and transients on time-scales ranging down to $\\sim 100$ d at levels ranging up to $\\sim 10^{-7}$ photon/(cm$^{2}$ s), outside the Galactic plane, both numbers being about a factor 10 higher in the plane. Variable emission detected from the known extragalactic source QSO 2206+650 is presumably not related to ETI activity. These results, limited though they are, are the first ever obtained in this field. They exclude the presence of \"human-scale\" antimatter-powered space probes (such as might be constructed by humans in this century [11]) within a radius of $\\sim 10$ AU, and more ambitious human-like crewed interstellar craft out to several thousand AU. They will be greatly improved by future high energy $\\gamma$-ray missions such as GLAST." + }, + "0112/astro-ph0112459_arXiv.txt": { + "abstract": "{ We present the results of 1.3~mm observations of the Crab Nebula, performed with the MPIfR bolometer arrays at the IRAM 30-m telescope. The maps obtained, of unprecedented quality at these wavelengths, allow a direct comparison with high-resolution radio maps. Although the spatial structure of the Crab Nebula does not change much from radio to millimetre wavelengths, we have detected significant spatial variations of the spectral index between 20~cm and 1.3~mm. The main effect is a spectral flattening in the inner region, which can be hardly explained just in terms of the evolution of a single population of synchrotron emitting electrons. We propose instead that this is the result of the emergence of a second synchrotron component, that we have tried to extract from the data. Shape and size of this component resemble those of the Crab Nebula in X rays. However, while the more compact structure of the Crab Nebula in X rays is commonly regarded as an effect of synchrotron downgrading, it cannot be explained why a similar structure is present also at millimetre wavelengths, where the electron lifetimes far exceed the nebular age. Our data, combined with published upper limits on spatial variations of the radio spectral index, also imply a low-energy cutoff for the distribution of electrons responsible for this additional synchrotron component. Although no model has been developed so far to explain the details of this component, one may verify that the total number of the electrons responsible for it is in agreement with what predicted by the classical pulsar-wind models, which otherwise are known to fail in accounting for the number of radio emitting electrons. This numerical coincidence can give indications about the origin of this component. We have also detected a spectral steepening at millimetre wavelengths in some elongated regions, whose positions match those of radio synchrotron filaments. The steepening is taken as the indication that magnetic fields in synchrotron filaments are stronger than the average nebular field. ", + "introduction": "The Crab Nebula is the prototype of synchrotron nebulae powered by a spinning-down pulsar, also known under the name of ``plerions'' (Weiler \\& Panagia \\cite{wei78}). This is an extensively studied object, and a wealth of information on the synchrotron nebula comes from detailed observations performed in various spectral ranges, like in radio, infrared, optical, UV and X rays. Modelling all the available data in a comprehensive frame represents a formidable task for the theory. Classical approaches to the modelling of the Crab synchrotron emission, like Pacini \\& Salvati (\\cite{ps73}) and Kennel \\& Coroniti (\\cite{kc84a}, \\cite{kc84b}), got some success. But more elaborate models can hardly get any substantial improvement with respect to the original approaches, partly because the geometric structure of the Crab Nebula is very complex, but probably also because the processes involved are not fully understood. When more quantitative and detailed modelling will be possible with other plerions we expect to face similar problems: in these respects a large part of the results on the Crab Nebula are likely to be exported to other objects. Considering just the total luminosity spectrum, Pacini \\& Salvati (\\cite{ps73}) successfully reproduced it from radio to optical, by simply assuming a pure power-law distribution for the injected ``particles'' (hereafter used to indicate relativistic electrons, as well as positrons): but in order to explain by their model the further spectral steepening in the X-ray range and beyond, an {\\it ad hoc} spectrum for the injected particles is required. On the other hand Kennel \\& Coroniti (\\cite{kc84b}) successfully reproduced the spectrum from optical to gamma rays, just assuming a power-law distribution (over a range of energies) for the particles accelerated at the termination shock of the pulsar relativistic wind. However their model fails in explaining the observed radio emission: the problem is that the best-fit wind model implies also an estimate of the total number of radio emitting particles injected into the nebula, which is at least a factor 100 lower than what measured. Up to now this discrepancy has been cured only by introducing some {\\it ad hoc} assumptions (see e.g.\\ Atoyan \\cite{ato99}). As far as the spatial structure of the nebula is concerned, it is not difficult to explain qualitatively its behaviour with frequency, namely the shrinking of the nebular size with increasing frequency: in fact the latter corresponds to increasing particle energy, and therefore decreasing synchrotron lifetimes. However quantitative approaches fail to reproduce the observed profiles, both in the Kennel \\& Coroniti (\\cite{kc84b}) and in the Pacini \\& Salvati (\\cite{ps73}) frameworks: the implications of the assumptions in the latter paper on the nebular spatial extent have been investigated by Amato et al.\\ (\\cite{aea00}). A common characteristic of the above models is that the particles are advected outwards with the magnetic field, following the MHD equations. Somehow better results are for instance obtained by including also diffusive processes, but only when an {\\it ad hoc} diffusion coefficient is taken. An alternative to the above scenarios relies on assuming the coexistence of two (or more) components of injected particles, with different spectra as well as with different spatial locations. But one may be unwilling to increase the complexity of the models, unless a stringent evidence in that sense comes out from the observations. Since adiabatic losses preserve the slope of the particles distribution, the most direct test on the presence of multiple components of the injected particles consists into measuring spatial variations of the synchrotron spectral index that cannot in any way result from a synchrotron downgrading (i.e.\\ the spectral softening consequent to synchrotron evolution of the emitting particles). This can be done observing at frequencies so low that the related particles are subject to negligible synchrotron losses, and therefore whose distribution retains the slope which had at the injection. Beforehand this kind of test had been done only at radio wavelengths (below with the term ``radio wavelengths'' we roughly indicate wavelengths above 1~cm). Previous claims of spatial variations of the radio spectral index (Velusamy et al.\\ \\cite{vea92}) have been then contradicted (Bietenholz et al.\\ \\cite{bea97}). Possibly some variations of the spectral index are present in the very central region, with scales of a few arcsec and associated to the ``wisps'' structures (Bietenholz \\& Kronberg \\cite{bk92}): however such result could be a mere artifact, originated by comparing data taken at different epochs, in the presence of rapidly moving wavy structures (Bietenholz et al.\\ \\cite{bea01}). In fact recent radio observations (Bietenholz et al.\\ \\cite{bea97}) strongly support the idea of a single injected distribution, by putting a tight upper limit, 0.01, to spatial variations of the spectral index, at least on scales larger than 16\\arcsec. In this paper we will show that millimetric wavelengths represent the most appropriate spectral range to investigate this issue in the Crab Nebula, by providing new pieces of information with respect to the radio. The map presented in this paper, with a 10.5\\arcsec\\ resolution, is by far better than the only map of the Crab Nebula previously published at these wavelengths (Mezger et al.\\ \\cite{mea86}, with only 120\\arcsec\\ resolution). The paper is organized as follows: in Sect.~2 we report on the observation parameters and on the data reduction; Sect.~3 describes the procedure by which our 1.3~mm map has been compared with a 20~cm radio map; features coming out from this comparison, namely the emergence of a second component in the inner regions and a general bending in filaments spectra, are respectively discussed in Sect.~4 and Sect.~5; Sect.~6 shows that the nature of the emission from the inner component is synchrotron; the morphology of the new synchrotron component is compared in Sect.~7 with maps at other wavelengths; in Sect.~8 we comment on possible spurious effects on our results deriving from time variability of the source; Sect.~9 concludes. ", + "conclusions": "We have measured inhomogeneities in the spectral index between radio and mm wavelengths, and we have shown that they could be better explained in terms of the emergence of a further synchrotron component, undetected at radio wavelengths, which is located in the inner part of the nebula. The transition in size and shape of the Crab Nebula, moving from the radio to the X-ray spectral range, is qualitatively explained in terms of synchrotron downgrading. But the fact that size and shape of component {\\bf B} at mm wavelengths resemble those seen in X rays cannot be explained in that way, since mm emitting particles are subject to only minor synchrotron losses. We suggest instead that two different synchrotron components coexist in the Crab Nebula and that the morphological transition taking place from radio to X rays requires a change of the relative importance of the two components. In order to fit the data, the energy distribution of particles emitting in component {\\bf B} requires to have a low-energy cutoff. With this, the total number of particles in component {\\bf B} are in agreement with what predicted by Kennel \\& Coroniti (\\cite{kc84b}) model. Finally, in synchrotron filaments we have found the evidence for a spectral break at a frequency lower than that averaged over the whole nebula. Although there may be some effects related to particle diffusion through filaments, we take a magnetic field in filaments higher than in their surroundings as the main cause of this effect." + }, + "0112/astro-ph0112345_arXiv.txt": { + "abstract": "We analytically investigate the formation of an HII region in the accreting envelope of a newborn star. Special care is taken to examine the role of ionizing radiation force. This effect modifies velocity and density distributions and thereby affects the expansion of the HII region. As a result, the upper limit of the stellar mass imposed by the growth of an HII region around a forming star is increased by a larger factor than the previous estimate. In particular, for a star forming out of metal-free gas, this mechanism does not impose a firm upper limit on its mass. ", + "introduction": "Stars are supposed to be born in the dense regions of gas clouds. The theory of the gravitational fragmentation of the gas clouds describes how these dense regions are formed inside the parental gas clouds. The typical mass scale of the dense region can be predicted by the Jeans mass of the cloud at the time of the fragmentation. However, a perturbation with a mass scale larger than this Jeans mass is always unstable. Therefore a dense region that is more massive than the typical Jeans mass can form depending on the initial fluctuation spectrum in the cloud. In this sense, the upper limit for the mass of stars cannot be obtained from the analysis of the gravitational fragmentation. Once, the origin of the observed upper limit of stellar mass (around $100M_{\\sun}$) was attributed to the instability of massive stars due to the $\\varepsilon$-mechanism (e.g., Schwarzshild \\& H\\\"{a}rm 1959). Since later studies revealed that this mechanism leads to moderate mass-loss rather than disruption of the stars (e.g., Appenzeller 1970), formation process, instead of the stability of massive stars, has been believed to limit the upper bound for the stellar mass (e.g., Nakano, Hasegawa, \\& Norman 1995). The star formation process is an accretion of ambient matter onto a protostar (stellar core) forming inside a protostellar cloud. In this scenario, the final mass of a star depends on how much mass the protostar can acquire. Although it is not yet clear what mechanism stops the accretion, increasingly strong feedback from a massive forming star is likely to halt the accretion. Larson \\& Starrfield (1971) were first to propose that a firm upper limit on the stellar mass is provided by the formation of an HII region around a forming star, as well as by the radiation force acting on dust grains in the accreting envelope. The formation of an HII region prevents further accretion in the following way: when the HII region reaches outer layers of a protostellar cloud, the temperature and then the pressure support surge by some large factor, so that the further infall of material is immediately halted. Larson \\& Starrfield (1971) concluded that both of those mechanisms set an upper limit of about 50 M$_{\\sun}$ on the masses of Pop I stars forming out of the present interstellar medium. The radiation force onto dust grains was found by subsequent studies to set more stringent mass upper limit (Appenzeller \\& Tscharnuter 1974; Kahn 1974; Yorke \\& Kr\\\"{u}gel 1977; Wolfire \\& Cassinelli 1987). Among them, Wolfire \\& Cassinelli (1987) suggested that even stars as massive as about 10 M$_{\\sun}$ cannot form by accretion owing to this mechanism unless dust is significantly depleted. However, for so-called Pop III stars forming out of metal-free gas, this mechanism of the radiation force does not work because of the absence of dust grains. The $\\varepsilon$-mechanism drives even milder mass-loss for those stars than Pop I massive stars (Baraffe, Heger, \\& Woosley 2001). According to recent simulations of the fragmentation of primordial clouds, the mass scale of fragments (protostellar clouds) is as large as $10^{3} M_{\\sun}$ (Bromm, Coppi, \\& Larson 1999; Abel, Bryan, \\& Norman 2000; Tsuribe 2001). Therefore, the growth of the HII region plays a crucial role in limiting the maximum mass of a Pop III star, if the upper limit by this mechanism is less than the mass scale of fragments. This motivates us to study here the upper mass limit of Pop III stars due to the formation of an HII region. In discussing the formation of an HII region, often the free-fall assumption has been imposed on the flow in the accreting envelope (e.g., Yorke 1986). The role of the momentum transfer to the gas due to ionizing radiation has been neglected. In the context of galaxy formation, however, its significance in dynamical and thermal evolution of the intergalactic medium has been pointed out by Haehnelt (1995). He showed, for the collapse of subgalactic clouds of $\\la 10^{10} M_{\\sun}$, even a radiation-pressure-driven bounce is possible. Taking this effect into account, in this paper, we study the formation of an HII region in the accreting envelope of a forming massive star. Although this effect is negligible as long as the radius of the HII region is small, it will turn out that inclusion of the ionizing radiation force alters the later evolution of the HII region. In the course of the expansion of the HII region, the flow becomes slower than the free-fall rate owing to the ionizing radiation force. The growth of the HII region is strongly suppressed by this effect. Consequently, the mass upper limit imposed by the expansion of the HII region is increased by a large factor. In particular, the formation of an HII region does not set any firm mass upper limit for Pop III stars. In Section 2, we briefly summarize the relevant aspects of the formation of an HII region in a free-falling envelope around a newborn star. The effect of ionization radiation force is included in Section 3. Finally, we provide a summary and discussion in Section 4. ", + "conclusions": "To estimate the mass upper limit of metal-free stars imposed by the formation of an HII region, we have found solutions for spherically symmetric and steady accretion flow onto a star emitting ionizing radiation. The behavior of the flow is determined by the velocity of the flow entering the HII region. If the flow velocity at the edge of the HII region goes beyond (below) a critical value, the flow is always accelerating (decelerating, respectively) in the HII region. When this value is equal to the critical value, the velocity remains constant in the HII region. In this critical flow, the radiation force due to ionizing photons exactly balances the gravity. We applied those solutions for a compact HII region forming around an accreting star. As the HII region grows in radius, the accreting flow evolves from an accelerating solution to a critical one. Soon after the critical flow is reached, accretion is halted by radiation force due to the ionizing radiation. However, even stars as massive as $10^{3}M_{\\sun}$ are unable to emit ionizing photons sufficient to halt the accretion. The halting by gas pressure is even more difficult. Therefore, contrary to the previous expectation, the formation of the HII region does not impose a stringent mass upper limit (at least up to $10^{3}M_{\\sun}$) on metal-free stars. We have estimated the upper limit of the mass of Pop III stars in relation to the formation of an HII region. Note that other effects (e.g., stellar wind, mass outflow due to pulsational instability, etc.) could grow in importance and might decrease the upper mass limit of Pop III stars below the value obtained in this paper. We leave detailed studies on these topics for future work. Also, in this paper, we have adopted very simplified assumptions, for example spherical symmetry, steady accretion, etc. Here, we discuss other complexities and possible deviations from our picture. We have considered only photoionization and electron scattering as sources of radiation force. Here, we mention briefly the radiation force due to Ly$\\alpha$ photons. Ly$\\alpha$ photons are emitted from recombination in the HII region. The flow inside the HII region is exerted by the radiation force due to Ly$\\alpha$ photons. We can extend our theory easily to include the Ly$\\alpha$ pressure , since, in the HII region, the radiation force due to the Ly$\\alpha$ photons is proportional to that due to ionization photons (equation 12 of Haehnelt 1995; see also Braun \\& Dekel 1989; Bithell 1990). However, according to Haehnelt(1995), the radiation force due to Ly$\\alpha$ photons is at most of marginal importance relative to that due to the ionizing photons. Thus, for the sake of simplicity, we have not included it. The HII region is surrounded by an HI layer, which is very optically thick to Ly$\\alpha$ photons. Without dust grains, which absorb Ly$\\alpha$ photons and reemit them as infrared photons, the Ly$\\alpha$ photons must diffuse out through the HI layer. In the course of this, the HI layer is pushed outward by those photons. Doroshkevich \\& Kolesnik (1976) argued that this mechanism expels the HI layer soon after the HII region is formed and thereby limits the mass of stars below 10 M$_{\\sun}$. However, Harrington (1973) showed that, even without dust, the two-photon emission process decreases the Ly$\\alpha$ photon density drastically, and the Ly$\\alpha$ radiation force is not dynamically important in an HI layer surrounding an HII region. In our case, supersonic motion in the accreting envelope also decreases the Ly$\\alpha$ photon density. Considering these facts, it is likely that the Ly$\\alpha$ radiation force does not play a significant role in our case. Therefore, we chose to neglect it here and assumed the free-fall outside the HII region. We have identified the base of an HII region as the stellar surface, and have taken the typical value of the stellar radius $10R_{\\sun}$ as the inner boundary radius $R_{\\rm in}$ of the HII region. However, when the accretion rate is high, a photosphere may be formed in the accreting flow (Wolfire \\& Cassinelli 1986; Stahler, Palla, \\& Salpeter 1986). In this case, $R_{\\rm in}$ should be taken as the radius of the photosphere, since we assume the HII region is optically thin to continuum absorption except for photoionization outside $R_{\\rm in}$. In spite of higher $R_{\\rm in}$, this effect generally works towards the smaller HII region, since $Q$ drops (Wolfire \\& Cassinelli 1986). Hence, our conclusion of higher upper mass limit than the former estimate remains the same. The spherical symmetry is clearly an oversimplification if the accreting matter has large angular momentum and an accretion disk is formed. Even in this case, the spherical symmetric flow is a good approximation outside the centrifugal radius. Suppose here that inside the centrifugal radius, there is a ``cavity'' around the disk. In this case, $R_{\\rm in}$ should be taken as the radius of the cavity. If $Q$ is unattenuated inside the cavity, the halting of the flow outside becomes easier because of higher $R_{\\rm in}$ and then lower $Q_{\\rm crit}$ (see eq. [\\ref{eq:qc_0}]). The accretion through the disk might continue, however (Nakano 1989). Those issues are still too speculative and beyond the scope of this paper. Although the accelerating solution is Rayleigh-Taylor stable, it might be unstable to % perturbations: if a portion of the flow becomes slightly slower than the rest, it becomes slower and slower relative to the average flow owing to the increased radiation force. In this case, density irregularities or blobs could be formed in the flow, and our spherically symmetric solution might be regarded as an approximate description of the average flow. Further study of the stability and dynamical evolution of the flow will be interesting." + }, + "0112/astro-ph0112035_arXiv.txt": { + "abstract": "{This paper is concerned with small angular scale experiments for the observation of cosmic microwave background anisotropies. In the absence of beam, the effects of partial coverage and pixelisation are disentangled and analyzed (using simulations). Then, appropriate maps involving the CMB signal plus the synchrotron and dust emissions from the Milky Way are simulated, and an asymmetric beam --which turns following different strategies-- is used to smooth the simulated maps. An associated circular beam is defined to estimate the deviations in the angular power spectrum produced by beam asymmetry without rotation and, afterwards, the deviations due to beam rotation are calculated. For a certain large coverage, the deviations due to pure asymmetry and asymmetry plus rotation appear to be very systematic (very similar in each simulation). Possible applications of the main results of this paper to data analysis in large coverage experiments --as PLANCK-- are outlined. ", + "introduction": "Many experiments are being designed for the observation of the Cosmic Microwave Background (CMB) anisotropies. From the maps of a given experiment operating with a non-circular (asymmetric) rotating beam, a certain angular power spectrum ($C_{\\ell}$ quantities) can be extracted. Different rotations can lead to distinct $C_{\\ell}$ coefficients and, the question is: How different are these coefficients? In other words, how relevant is the effect of the rotation strategy on the resulting angular power spectrum? In a previous paper (Arnau \\& S\\'aez 2000), it was shown that, in the absence of rotation and when the level of instrumental noise is low enough, the effect of a non-circular beam can be subtracted --namely, the beam can be deconvolved-- using the Fourier transform. This subtraction can be performed in such a way that the resulting spectrum, after deconvolution, is very similar to the true one. That is possible if the number of pixels inside the beam, $N_{\\mbox{in}}$, is not too great. Indeed, $N_{\\mbox{in}}$ cannot be much greater than 10; however, if the beam rotates, the deconvolution is not possible. Nobody has described either the importance of beam rotation or a method to eliminate its effects. The main goal of this paper is the estimation of the effects due to rotation. In Arnau \\& S\\'aez 2000 (and also in S\\'aez, Holtmann \\& Smoot, 1996, and S\\'aez\\& Arnau, 1997). a sort of modified angular power spectrum was used. Here, we extract the standard $C_{\\ell}$ quantities from a certain number of squared patches of the sky. Recently, Wu et al. (2001) have proposed a method for data analysis in the case of asymmetric beams. This method is based on an optimal circular beam associated to the asymmetric one. The effects of beam rotation are not studied at all by these authors. Although our methods apply to CMB anisotropy experiments in general, we will pay particular attention to PLANCK mission (scheduled by ESA for 2007). As it was emphasized in Burigana et al. (1998), {\\em beam responses are typically nonsymmetric for detectors de-centred from the telescope focus}. Taking into account that CMB anisotropy experiments require observations at different frequencies, various detectors are necessary, which must be distributed as close as possible from the focus; for instance, in the PLANCK mission, around one hundred of detector (bolometers and radiometers) must be distributed in the focal plane. If the focal plane rotates (rotation of the telescope around the spin axis), the beams do. The effect of this rotation deserves attention. Furthermore, there are various identical detectors for each frequency, which are located at different positions in the focal plane and, consequently, the deformations of these beams would be different (identical) if they are located at different (the same) distances from the optical focus; nevertheless, even for identical deformations, the orientations of the resulting asymmetric beams would be different. The motion of the line of sight through the sky also produces a beam asymmetry. The effective beam diameter $\\theta_{_{\\mbox{FWHM}}}$ appears to be enhanced in the direction of this motion (see Hanany, Jaffe \\& Scannapieco 1998). This small effect is due to the beam displacement during the measurement process. It is not taken into account in this paper. Beam rotation depends on the particular experiment under consideration. Given a pixelisation, the beam centre points towards a given pixel a certain number of times, $N_{\\mbox{p}}$, and, then, the temperature assigned to this pixel is an average of the temperatures corresponding to each of the $N_{\\mbox{p}}$ measurements. The fact that measurements from various beam orientations are averaged could be important. In the case of PLANCK mission, a rough estimate of number $N_{\\mbox{p}}$ is given in S\\'aez \\& Arnau (2000). Here, it is worthwhile to improve a little on that estimate. The satellite has been designed in such a way that: (i) it will cover the full sky in seven months, with a coverage which can be considered as uniform in most part of the sky, (ii) its line of sight will move around a big circle on the sky completing a turn each minute and, (iii) it will move around the same circle for two hours (120 turns). On account of these facts, if the pixel size is $\\Delta$ and the angle subtended by the motion of the line of sight between two successive measurements is $\\Delta \\alpha = \\zeta \\theta_{_{\\mbox{FWHM}}}$, where $\\theta_{_{\\mbox{FWHM}}}$ is the beam diameter, then, the average number of measurements per pixel (in a seven months observing period) is $N_{\\mbox{p}} = 42 \\Delta^{2} / \\zeta \\theta_{_{\\mbox{FWHM}}}$, where all the angles are given in arc-minutes (see S\\'aez \\& Arnau, 2000, for comparison) and, furthermore, the average number of measurements per pixel performed while the line of sight turns 120 times around a given circle is $N_{\\mbox{pc}} = 120 \\Delta / \\zeta \\theta_{_{\\mbox{FWHM}}}$. From these formulae, it follows that the average number of circles passing by a pixel --during seven months-- is $N_{\\mbox{c}} = N_{\\mbox{p}}/N_{\\mbox{pc}} \\simeq \\Delta/3$, this result is consistent with the fact that, for a given observational strategy, the number $N_{\\mbox{c}}$ is expected to be dependent only on the pixelisation. Of course, it is independent on beam asymmetry. The number of measurements corresponding to different orientations could be important for the effect we are looking for, which is produced by the rotation of asymmetric beams. The larger the pixel size, the better the situation (the greater $N_{\\mbox{c}}$). Since the detectors are rigidly attached to the focal plane, any beam has almost the same orientation each time it crosses a given pixel during its motion (120 turns) along a given circle; however, this orientation changes from circle to circle. From the above comments and estimates, it follows that the average number of measurements per pixel corresponding to different beam orientations is $N_{\\mbox{c}}$. If the full sky is covered two times and, the second coverage is not identical to the first one, this average number would be $2 \\Delta/3$. For $5^{\\prime } < \\Delta < 10^{\\prime}$, this number ranges from 3.3 to 6.6. Nevertheless, there are various detectors in the focal plane for each frequency and, by assuming that all the beams have the same shape but different orientations, the above $N_{\\mbox{c}}$ number can be multiplied by the number of beams. ", + "conclusions": "In the absence of beam, the pixelisation effect and the REPC have been disentangled to conclude that pixelisation produces very systematic deviations with respect to the true angular power spectrum. This conclusion has been obtained using two very different methods for simulations and data analysis (see Sect. 3) We have studied the deviations in the angular power spectrum produced by the rotation of an asymmetric beam. Two rotation strategies have been considered. One of them (SR) is similar to that of future experiments as PLANCK. The second strategy (RR) is very different from SR, and it has been introduced for comparisons. Maps with and without the dust and synchrotron radiations from the Milky Way (at $100 \\ GHz$) have been considered. In Sect. 4, the rotation effects corresponding to different cases have been described and compared, now let us present some general comments. If radiation from the galaxy is not considered, the most important conclusion is that rotation effects are very systematic for any rotation strategy and $f_{\\mbox{sky}} = 0.39$. They are so systematic that we can subtract the deviations appeared in a 50-simulation, from the spectrum of another one, to recover very well the spectrum corresponding to the nonrotating beam (except for small deviations which seem to be essentially due to the REPC). Furthermore, the resulting effects depend on the rotation strategy, in particular, for large $\\ell$ values and, consequently, they must be estimated --using simulations-- in each particular case. Radiation from the galaxy --which can be seen as a non homogeneous an non isotropic statistical field-- contributes significantly to the observable signal, except in the case of the polar galactic regions (G5--G6). In the G1--G2 and G3--G4 cases, the effect of beam rotation is significantly, but not dramatically, different from that obtained in the absence of foregrounds. After the deviations in the angular power spectrum due to beam asymmetry and rotation have been estimated and characterized (the main goal of this paper) and, after proving that beam effects are very systematic, some practical applications can be easily outlined. Take the CMB power spectrum corresponding to a certain theoretical model of structure formation in a given universe, take also a model for the foregrounds, a pixelisation, the asymmetric beam for a given frequency, and the rotation strategy of an experiment with a large enough coverage (i.e. PLANCK), and then, use a simulation --as the 50-simulations of this paper or similar-- to find the spectrum $C_{\\ell}^{^{\\mbox{CMB}}}$ after smoothing with the asymmetric rotating beam. Repeat the simulation a large enough number of times and verify that the resulting $C_{\\ell}^{^{\\mbox{CMB}}}$ spectra are similar in all cases (systematic character). Finally, use the deviations among the resulting spectra to assign an error bar to $\\langle C_{\\ell}^{^{\\mbox{CMB}}} \\rangle$. Use these data --obtained from simulations-- to answer the following question: Is the theoretical model under consideration compatible with the observational data from the experiment? In order to find the answer, the observational data could be analyzed as follows: (i) Eliminate a part of the instrumental noise using an appropriate method (wavelets, Fourier transform, and so on), (ii) Separate components (CMB, synchrotron from our galaxy, and so on) taking into account the frequency dependences, but keeping beam smoothings unaltered (usually, the beams are eliminated at this stage under simplifying assumptions and without considering rotation), (iii) use the map of the CMB component --which has already been separated from foregrounds-- to extract the experimental spectra, $C_{\\ell}^{^{\\mbox{CMB}}}(exp)$, and finally (iv) compare $C_{\\ell}^{^{\\mbox{CMB}}}$ with $C_{\\ell}^{^{\\mbox{CMB}}}(exp)$ and study if these spectra can be identified taking into account the error bars. If they can, the theoretical model is compatible with observations. Note that --at the last step of the process-- we compare a simulated spectrum with an observational one, and note also that both spectra are obtained from maps which have been smoothed with the same rotating asymmetric beam; hence, the proposed method for data analysis includes beam rotation, treating it (after verification) as the source of a very systematic effect. Of course, this method has been only outlined, and much more work would be necessary before implementation." + }, + "0112/hep-th0112128_arXiv.txt": { + "abstract": "By considering a simplified but exact model for realizing the ekpyrotic scenario, we clarify various assumptions that have been used in the literature. In particular, we discuss the new ekpyrotic prescription for passing the perturbations through the singularity which we show to provide a spectrum depending on a non physical normalization function. We also show that this prescription does not reproduce the exact result for a sharp transition. Then, more generally, we demonstrate that, in the only case where a bounce can be obtained in Einstein General Relativity without facing singularities and/or violation of the standard energy conditions, the bounce cannot be made arbitrarily short. This contrasts with the standard (inflationary) situation where the transition between two eras with different values of the equation of state can be considered as instantaneous. We then argue that the usually conserved quantities are not constant on a typical bounce time scale. Finally, we also examine the case of a test scalar field (or gravitational waves) where similar results are obtained. We conclude that the full dynamical equations of the underlying theory should be solved in a non singular case before any conclusion can be drawn. ", + "introduction": "Modern ideas of particle physics, such as superstring, $M-$theory~\\cite{sugrastring} or quantum gravity~\\cite{QG}, cannot in general be subject to experimental constraints because of the enormous energies (usually of the order of the Planck mass) at which they are supposed to become effective. According to recent theoretical developments~\\cite{largeD,RS}, there is hope that space possesses more than three large dimensions and that these extra dimensions might turn out to be observable in a not too distant future. The majority of the theoretical models that have been built so far are however based on extremely high energy extensions of the standard particle physics model, and thus currently need to be tested by the yardstick of cosmology, the latter being the only playground at which those theories could have acted. According to the now standard paradigm that describes the early universe and that is expected to stem from such high energy particle models, a phase of superluminal accelerated expansion known as inflation~\\cite{inflation} preceded the radiation-dominated epoch. Up to now, no model has come as close to being a reasonable challenger to solve the standard cosmological puzzles (flatness, homogeneity and monopole excess). The extra bonus provided by the inflationary phase is that it leads naturally to a scale-invariant density fluctuation spectrum that seems to be in agreement with the observations. Inspired by the recent developments of $M-$theory~\\cite{heterotic}, in particular through Ref.~\\cite{HW}, and invoking brane cosmology, recent work~\\cite{ekp,ekpN,perturbekp} claimed to be able to solve all the aforementioned problems as well, including a new way of producing primordial cosmological perturbations. Although the model, both in its ``old''~\\cite{ekp} and ``new''~\\cite{ekpN} versions is plagued with many difficulties~\\cite{pyro}, as a potential alternative to inflation (see also Ref.~\\cite{PBB} in that respect), it is worth examining in detail, would it be only to re-enforce the confidence we may have in the latter. In both the original and most recent versions, the universe is supposed to consist of a four dimensional (visible) brane evolving in a higher (in practice 5) dimensional bulk. By assuming the brane to be a Bogomolnyi-Prasad-Sommerfield (BPS) state~\\cite{BPS}, one ensures that the curvature $\\Ka$ vanishes, thus addressing the flatness problem. To begin with, another brane, that can be either a light bulk brane~\\cite{ekp}, or the other (hidden) boundary brane~\\cite{ekpN}, moves freely in the bulk until it collides with the visible brane. The collision time is interpreted as the hot big bang at which point the model is made to match the standard cosmological model. Apart from the collision time, the theory, which can be seen as effectively four dimensional in the long wavelength limit, relies on the General Relativity (GR) theory together with some extra fields. In this effective 4D model, the Universe collapses, experiences a bounce at some instant in time, and starts expanding. As far as cosmological perturbations are concerned, only GR calculations have been discussed up to now. The pre-impact phase has been the subject of many tentative calculations of the perturbation spectrum that would be generated by quantum perturbations of the brane~\\cite{perturbekp}. A general agreement has now been reached~\\cite{perturbekp,Lyth,BF,Hwang} that the curvature perturbation spectrum $P_{\\zeta }$ has spectral index $n_\\zeta=3$, while that of the Bardeen potential $P_{\\rm \\Phi }$ ends up with $\\nS=1$, i.e., a scale invariant spectrum. On the other hand, the spectra of $\\Phi$ and $\\zeta$ are identical in the post-impact phase, and enter the Cosmic Microwave Background Radiation (CMBR) multipole moments. It is therefore of utmost importance to obtain full knowledge of these spectra not only in the pre-impact phase, but also after the bounce has occurred, i.e., at times that are observable now. In other words, the fate of $\\Phi$ and $\\zeta$ through the bounce is the main issue before any conclusion regarding the model can be drawn. Only a few definite statements can be done about the bounce epoch. The first, which was advocated by many authors, is that GR does hold during it, or, stated differently, that it lasts sufficiently little that corrections to GR can be regarded as negligible. Lacking the actual theory, this is the only statement that can be endowed with a predictive power. To begin with, it implies that there was no singularity, and, if the null energy condition is to be satisfied, that space is positively curve, i.e., $\\Ka=1$. Under these conditions, ordinary perturbation theory~\\cite{Bardeen,perturb} can be applied. By assuming continuity of the Bardeen potential and the well known conserved quantity $\\zeta$ (defined below), it was then found~\\cite{BF} that the scale invariant spectrum does not survive the bounce, with the actual resulting spectrum being much lower than the observed one. The temporary conclusion of this fact is that in order that the ekpyrotic model be still compatible with the observation, a new procedure must be applied to the bounce. Arguing against GR during the bounce epoch sounds natural, as in particular either the real theory is at least 5-dimensional, or, worst indeed, in the case of the new scenario~\\cite{ekpN}, the manifold becomes (curvature) singular there, obviously leading to a breakdown of ordinary GR across the bounce. In this case, a new criterion should be derived to replace the ordinary junction conditions. Such a criterion was provided in Ref.~\\cite{perturbekp}, although without a physical motivation, leading to the recovery of the observationally correct spectrum. The very exhibition of junction conditions leading to a scale invariant spectrum could then be seen as a hint that constructing a realistic theory satisfying observational constraints was not impossible. Even if one is prepared to accept such drastic changes in the standard cosmological picture, one might wonder as to the use of perturbation theory on top of an otherwise singular background~\\cite{Lyth}. Moreover, it should be mentioned that although the old scenario, because describable as an effective bounce occurring at a low enough temperature, was avoiding the over-production of grand unified scales monopoles~\\cite{monop}, the new model, being singular, poses this problem in a way which is as acute as it was before the advent of inflation. Finally, the puzzle of trans-Planckian scales~\\cite{transPl}, quoted in Ref.~\\cite{reply} as a caveat for inflation, can be transposed in the new ekpyrotic model in the same words. This article is organized as follows. After a brief reminder of the ekpyrotic model of the universe (Sec.~\\ref{sec:ekp}), we examine in detail the junction conditions suggested in Ref.~\\cite{perturbekp} (Sec.~\\ref{sec:pert}). We concentrate in particular on the fact that this proposed criterion rests on an altogether arbitrary (hence unphysical) normalization function, so that whatever spectrum can be obtained: obtaining a scale invariant spectrum in this model thus turns out to be equivalent to imposing it from the outset. We also demonstrate that the new prescription leads to an incorrect prediction in the exact case of a radiation to matter domination transition. We then consider a second possibility, i.e., we examine an effective bounce in a context where the linear perturbation theory is still valid. We therefore considered first, in section~\\ref{sec:hydro}, the simplest case in which not only does GR apply, but also in which all the calculations can be performed analytically and consistently (indeed providing a nice textbook example for cosmological perturbation theory illustration), namely that of a ${\\cal K}=1$ bouncing universe with hydrodynamic perturbations~\\cite{ppnpn1}. Then, using this toy model, we examine how the relevant perturbed quantities behave through the bounce. We pay special attention to the ``short time bounce limit'' (this is related to the question ``how sharp is sharp'' evoked in Ref.~\\cite{Lyth}) and study whether, in this limit, the bounce can be considered as a surface where the equation of state jumps. If so this would allow us to use the standard junction conditions. The second example that one can treat completely is by considering a test scalar field. Indeed, in this case, one does not need to specify what the origin of the background evolution is. In section~\\ref{sec:scalar}, we calculate the spectrum of a spectator scalar field in such a bouncing background. Assuming no strong deviation from GR at the perturbed level (we remind that such deviations are necessary in the bounce region), and $\\Ka=0$, this also gives the tensor perturbation spectrum. The description of a bouncing universe with ${\\cal K}=0$ requires special care, as GR does not allow for such a configuration to take place unless the Null Energy Condition (NEC) is violated. Although this case is clearly contrived, it provides at least an example where some arguments presented recently in the literature can be implemented concretely, at the level of equations. ", + "conclusions": "The conclusions that can be drawn from this work is that it seems impossible to apply any known and well motivated criterion to pass through a bounce, whether regular or singular, in a model independent way as all quantities of interest explicitly depend on the details of the underlying model. The ekpyrotic model, although a potentially interesting alternative to the inflationary paradigm, does pass through such a bounce. Therefore, if one really wants to calculate the spectrum in the ekpyrotic universe then it seems necessary, first, to consider a situation where there is no divergence and, second, to provide us with the actual (maybe five-dimensional) equations of motion during the bounce, knowing that these equations cannot be those of GR." + }, + "0112/astro-ph0112565_arXiv.txt": { + "abstract": "In this paper we propose that the accelerating expansion of the present matter-dominated universe, as suggested by the recent distance measurements of type Ia supernovae, is generated along with the evolution of space in extra dimensions. The Einstein equations are first analyzed qualitatively and then solved numerically, so as to exhibit explicitly these patterns of the accelerating expansion in this scenario. A fine-tuning problem associated with such a scenario is also described and discussed. ", + "introduction": "The recent distance measurements of type Ia supernovae suggest an accelerating expansion of the present universe \\cite{Perlmutter:1999np,Riess:1998cb}. In many of the current cosmological models, the present accelerating expansion is driven by an energy source called ``dark energy'', with a positive cosmological constant \\cite{Lambda models} or the so-called ``quintessence'' (a slowly evolving scalar field \\cite{Caldwell:1998ii,ComplexQ}) as a possible candidate. Instead of attributing this acceleration to the mysterious dark energy, we consider in this paper the possible existence of extra spatial dimensions and explore the feasibility of producing the present accelerating expansion via the evolution of these extra dimensions. The application of extra dimensions is a general feature in theories beyond the standard model, especially in theories for unifying gravity and other forces, such as superstring theory. These extra dimensions should be ``hidden'' for consistency with observations. Various scenarios for ``hidden'' extra dimensions have been proposed, for example, a brane world with large compact extra dimensions in factorizable geometry proposed by Arkani-Hamed \\emph{et al.} \\cite{Arkani-Hamed} (see also \\cite{Antoniadis:1990ew}), and a brane world with noncompact extra dimensions in nonfactorizable geometry proposed by Randall and Sundrum \\cite{Randall&Sundrum}. In this paper, we employ the simplest scenario: small compact extra dimensions in factorizable geometry, as introduced in the Kaluza-Klein theories \\cite{Kaluza&Klein}. We study spatially homogeneous, isotropic, perfect-fluid cosmological models in $(1+3+n)$ dimensions where $n$ is the number of extra dimensions. In Sec.\\ \\ref{general features}, we first obtain, from the Einstein equations, some general features of the evolution of the higher-dimensional space, especially for a radiation-dominated universe and a (nonrelativistic) matter-dominated universe. In Sec.\\ \\ref{accel 3-space}, we then explore the possibility of producing an accelerating expansion of the ordinary three-space via the evolution of `extra space' for a matter-dominated universe. We analyze the Einstein equations to show qualitatively the behavior of this evolution and obtain numerical solutions which illustrate explicitly the accelerating expansion of the ordinary three-space along with the collapse of the extra space. We note that, while the Kaluza-Klein cosmology and inflation in higher dimensional space-time in connection with the early universe were discussed widely in the 1980s \\cite{Freund:1982pg,KK-cosmology} (for a review, see \\cite{Kolb:1990vq}; for recent work, see \\cite{Arkani-Hamed:1999gq,Mongan:2001cr}), the focus of this paper has to do with the present accelerating matter-dominated universe. ", + "conclusions": "We have investigated the scenario of producing the accelerating expansion of the present universe via evolving small extra dimensions. For a radiation-dominated universe, such as our early universe, we obtain a stable solution with static extra dimensions. Accordingly, the existence of extra dimensions may have no significant influence on the evolution of the ordinary three-space. This is a good feature which we need for preserving the concordance between observations and current theories regarding the early (radiation-dominated) universe, especially for primordial nucleosynthesis. On the contrary, such a solution with static extra dimensions does not exist for the present matter-dominated universe. The features of the evolution can also be read off from Eq.\\ (\\ref{alpha-volume relation}), or Eqs.\\ (\\ref{alpha-volume for RD}) and (\\ref{alpha-volume for MD}), which are derived from Eq.\\ (\\ref{alpha-volume relation}). Equation (\\ref{alpha-volume for RD}) shows the decreasing expansion rate of the extra space along with the increase of the $(3+n)$-dimensional volume $V_{3+n}$. This implies the stability of the solution with static extra dimensions in the radiation-dominated universe for the case of $k_a=k_b=0$ as already mentioned above. On the other hand, Eq.\\ (\\ref{alpha-volume for MD}) shows two possible evolution patterns of the matter-dominated universe: (i) The expansion rates of the ordinary three-space and the extra space tend to catch up with each other along with the increase of the (3+n)-dimensional volume $V_{3+n}$. (ii) One of these two expansion rates is positive and increasing, while the other is negative and decreasing, along with the decrease of the (3+n)-dimensional volume $V_{3+n}$. We note that an increasing positive expansion rate represents an accelerating expansion. A quantitative analysis of the matter-dominated case with $k_a=k_b=0$ leads to Fig.\\ \\ref{accel-decel plot}, which indicates four classes of evolution path. A universe that decelerates first and then accelerates is included in one of them. Therefore the accelerating expansion of the present universe may be appropriately described in this scenario. In addition, the case with two extra dimensions is analyzed in detail. The five resultant evolution paths we draw demonstrate the existence of a critical value for the initial condition $\\eta_0$, which divides two classes of path: the one in which the universe decelerates first and then accelerates and the other in which the universe always decelerates. We note that the critical value $\\eta_{cr}$ is exactly the parameter $K_{rep}$, a ``repeller'' in the flow diagram. However, the existence of the critical value (or the ``repeller'') also implies a fine-tuning problem, i.e., the initial value of $\\eta$ has to be chosen delicately so that it is close enough to the critical value $\\eta_{cr}$ in order to possess a long enough decelerating epoch followed by an accelerating epoch. The existence of extra dimensions is a general feature in theories beyond the standard model in particle physics. It may manifest itself as a source of energy in the ordinary three-space, such as ``effective'' dark energy or even ``effective'' dark matter. The geometrical structure and the evolution pattern of extra dimensions therefore may play an important role in cosmology. In this work we study a simple scenario of extra dimensions that is subject to a fine-tuning problem. Nevertheless, other scenarios with richer structures, such as those in \\cite{Arkani-Hamed} and \\cite{Randall&Sundrum}, may also provide suitable evolution patterns and are worthy of being further investigated. {\\em Note added}. For the sake of simplicity, we have in this paper considered only the case with $k_a=k_b=0$, i.e., both our ordinary three-space and the extra space are flat. The general situations with nonzero $k_a$ or $k_b$ clearly may offer many interesting possibilities and are currently under serious investigation." + }, + "0112/astro-ph0112086_arXiv.txt": { + "abstract": "We present new {\\it HST} I-band images of a sample of 77 nearby, late-type spiral galaxies with low inclination. The main purpose of this catalog is to study the frequency and properties of nuclear star clusters. In 59 galaxies of our sample, we have identified a distinct, compact (but resolved), and dominant source at or very close to their photocenter. In many cases, these clusters are the only prominent source within a few kpc from the galaxy nucleus. We present surface brightness profiles, derived from elliptical isophote fits, of all galaxies for which the fit was successful. We use the fitted isophotes at radii larger than $2\\as$ to check whether the location of the cluster coincides with the photocenter of the galaxy, and confirm that in nearly all cases, we are truly dealing with ``nuclear'' star clusters. From analytical fits to the surface brightness profiles, we derive the cluster luminosities after subtraction of the light contribution from the underlying galaxy disk and/or bulge. ", + "introduction": "Over the past decade, high dynamic range observations with modern CCD detectors have shown that compact stellar nuclei are a common feature of spiral galaxies of all Hubble types. For example, \\cite{mat97} found 10 objects with compact nuclear star clusters in a survey of 49 southern, very late-type spirals. However, as one progresses along the Hubble sequence towards earlier types, the increasingly luminous bulge component with its steeply rising surface brightness profile makes the identification of an additional, unresolved cluster extremely difficult. It therefore took the unique spatial resolution of the {\\it Hubble Space Telescope (HST)} to demonstrate that \\nc s are a common phenomenon also in earlier Hubble types \\citep*[e.g.][]{car98}. The {\\it HST} currently provides the only means to investigate the structural properties of nuclear star clusters, as demonstrated by \\cite{mat99}, and to cleanly separate their emission from the underlying galaxy disk/bulge. Despite the recent progress, the formation mechanism of nuclear star clusters remains largely a mystery. Intuitively, there are good reasons to expect matter accumulation in the deep gravitational wells of galaxies with massive bulges, and hence active star formation in their nuclei. In contrast, the gravitational force all but vanishes in the centers of pure disk galaxies with shallow surface brightness profiles and without any discernible bulge component. In these galaxies, the dynamical center is not a ``special'' place and it is far from obvious how a massive stellar cluster could have formed there. The shallow gravitational potential might provide a natural explanation for the fact that spirals of late Hubble type are not known to contain super-massive black holes. On the other hand, nuclear star clusters can be extremely compact: the nucleus of M33, for example, has likely undergone core collapse and is as compact as any known globular cluster \\citep{kor93}. So far, no satisfying explanation has been put forward to explain the high gas densities that must have been present in the nuclei of these shallow disk galaxies to enable the formation of such massive and compact objects. It is also unknown whether nuclear star clusters form repeatedly or only once - a question with important implications for the dynamical and morphological evolution of their host galaxies. To make progress along this line, it is essential to obtain the age distribution of nuclear star clusters. So far, reliable age estimates exist for only a handful of nuclear star clusters. Interestingly, most of them appear to be rather young: our Galaxy has a central stellar cluster with an age of only $\\sim 3$~Myrs \\citep{kra95}, and both M31 and M33 have blue nuclei that are very likely young star clusters \\citep{lau98}. More recently, we have published \\nc\\ ages derived from ground-based spectroscopy for IC~342 \\citep*[$\\rm \\leq 60\\>$Myrs,][]{boe99}, and NGC~4449 \\citep*[$6-10\\>$Myrs,][]{boe01,gel01}. In addition, the dominant stellar population of the nuclear cluster in NGC 3227 is less than 50~Myrs old \\citep*{sch01}. However, it is possible (and in fact quite likely) that ground-based observations predominantly target the brightest and hence youngest clusters. In order to get a more representative picture of nuclear star clusters, it is important to study a galaxy sample which is free from selection effects that favor the high end of the \\nc\\ luminosity range. In this paper, we describe the results of an {\\it HST} I-band imaging survey of an unbiased sample of nearby, face-on, very late-type spirals (Scd or later). The main goals of the survey are (a) to determine the frequency of nuclear star clusters in very late-type spirals, (b) to derive their luminosity and size distribution, (c) to compare their properties to those of nuclear star clusters in earlier Hubble types which have been more extensively studied with {\\it HST} \\citep{car97,car98,car01,car99}, and (d) to provide a source catalog for follow-up spectroscopic observations to age-date their stellar populations. The main purpose of this paper is to present the complete dataset. In a companion paper (B\\\"oker \\ea\\ 2002, in preparation), we describe the statistics of the full sample and investigate whether the properties of nuclear star clusters correlate in any way with those of their host galaxies. This paper is organized as follows: in \\S~\\ref{sec:obs}, we describe our sample selection criteria, the observational strategy, and the data reduction procedure, and we present the final images as well as the results of the isophotal analysis. In \\S~\\ref{sec:disc}, we discuss whether the clusters indeed occupy the nuclei of their host galaxies, and how they compare to other luminous star clusters observed in a variety of starburst environments. We conclude in \\S~\\ref{sec:concl}. ", + "conclusions": "\\label{sec:concl} We have presented a catalog of {\\it HST}/WFPC2 I-band images of an unbiased sample of 77 nearby, late-type spiral galaxies with low inclination. From an isophotal analysis of the images, we demonstrate that about 75\\% of the sample galaxies host a compact, luminous stellar cluster at or very close to their photocenter. These clusters often are completely isolated from other comparable structures, emphasizing that even in the relatively shallow potential wells of late-type galaxy disks, the center is well-defined, and has a unique star formation history. From analytical fits to the surface brightness profiles, we determine the flux attributable to the cluster. The distribution of absolute cluster luminosities has a FWHM of 4 magnitudes, and a median value of $\\rm M_I = -11.5$, comparable to young super star clusters in starbursting galaxies. Together with initial estimates of their size distribution, this suggests that \\nc s in spiral galaxies of the latest Hubble types are a fairly homogenous class of objects. The dataset is a representative survey of late-type spiral galaxies in the local universe, and as such yields a valuable source catalog for spectroscopic follow-up observations which are needed to further constrain the star formation history of \\nc s. We have begun such a follow-up program both with {\\it HST} and ground-based observatories." + }, + "0112/astro-ph0112423_arXiv.txt": { + "abstract": " ", + "introduction": "One of the most important quantities in theoretical models of AGNs is the black hole mass (M$_{BH}$) that, together with the total luminosity, defines the fraction of the Eddington luminosity at which the AGN is emitting. Determination of M$_{BH}$ in AGN is difficult mainly because of the bright emission from the nucleus and their large distance. The main method that has proved to be successful in AGN is reverberation mapping, which is extremely time consuming and gives results on M$_{BH}$ that depend on the assumed geometry of the accretion disk. Therefore, only for a few well studied quasars and Seyfert galaxies M$_{BH}$ is known (see e.g. [7], [11] and references therein). This method cannot obviously be employed for BL Lac objects because they lack prominent emission lines. Therefore other methods need to be applied to infer M$_{BH}$ for BL Lacs. The discovery of a relation between M$_{BH}$ and the luminosity of the bulge in nearby early-type galaxies offers now a new tool for estimating the mass of the central BH (see e.g. review [10]). This has been done for two samples of nearby quasars [9] and BL Lacs [12]. Recently, a tighter correlation was found relating M$_{BH}$ with the central stellar velocity dispersion $\\sigma$ of the spheroidal component in nearby galaxies [5,2], that can also be used to estimate M$_{BH}$ in AGN. The relationship appears to predict more accurately [10] M$_{BH}$, but requires the measurement of $\\sigma$ in the host galaxies of AGN that is difficult to obtain, in particular for objects at moderately high redshift and with very luminous nuclei. On the other hand, for BL Lacs that have relatively fainter nuclei than quasars, this measurement (at least for low redshift objects) can be secured with observations at medium-sized telescopes. We present here the first estimates of stellar velocity dispersion of BL Lacs from our ongoing program aimed specifically at deriving M$_{BH}$ from the M$_{BH}$ -- $\\sigma$ correlation. We selected a sample of nearby (z$<$0.2) BL Lacs for which high quality images were obtained either from the ground using the Nordic Optical Telescope (NOT) [3] or with HST+WFPC2 [13,4]. From the images a characterization of the host galaxies and of the nuclear luminosity are obtained. This allows us to compare M$_{BH}$ with the mass (and the luminosity) of the host galaxy and also to evaluate the Eddington ratio, provided that the nuclear emitted power is corrected for the beaming factor. Moreover, a comparison of M$_{BH}$ for BL Lacs with different jet/ disk luminosities can be used to test the hypothesis (see e.g. [8]) that the accretion rate changes from largely sub-Eddington, for low luminosity weak-lined sources, to near-Eddington for high luminosity, strong-lined sources. If the accretion rate in terms of Eddington ratio were the same in both classes, the BH masses should differ almost by three orders of magnitude. ", + "conclusions": "" + }, + "0112/astro-ph0112437_arXiv.txt": { + "abstract": "We study the dynamics of phase transitions in the interstellar medium by means of three-dimensional hydrodynamic numerical simulations. We use a realistic cooling function and generic nonequilibrium initial conditions to follow the formation history of a multiphase medium in detail in the absence of gravity. We outline a number of qualitatively distinct stages of this process, including a linear isobaric evolution, transition to an isochoric regime, formation of filaments and voids (also known as ``thermal'' pancakes), the development and decay of supersonic turbulence, an approach to pressure equilibrium, and final relaxation of the multiphase medium. We find that 1\\%-2\\% of the initial thermal energy is converted into gas motions in one cooling time. The velocity field then randomizes into turbulence that decays on a dynamical timescale $E_k\\propto t^{-\\alpha}$, $1\\lsim\\alpha\\lsim2$. While not all initial conditions yield a stable two-phase medium, we examine such a case in detail. We find that the two phases are well mixed with the cold clouds possessing a fine-grained structure near our numerical resolution limit. The amount of gas in the intermediate unstable phase roughly tracks the {\\em rms} turbulent Mach number, peaking at 25\\% when ${\\cal{M}}_{rms}\\sim8$, decreasing to 11\\% when ${\\cal{M}}_{rms}\\sim0.4$. ", + "introduction": "Thermal instability (TI) has many implications in astrophysics (e.g., a clumpy interstellar medium [ISM], stellar atmospheres, star formation, globular cluster and galaxy formation, etc.; see Meerson 1996 for a recent review). The instability may be driven by radiative cooling of optically thin gas (radiation-driven TI) or by exothermic nuclear reactions \\citep{schwarzschild.65}. Linear stability theory for a medium with volumetric sources and sinks of energy in thermal equilibrium was developed by \\citet{field65}, who identified three unstable modes: the {\\em isobaric} mode (the pressure-driven formation of condensations not involving gravitation) and the two {\\em isentropic} modes (the overstability of acoustic waves propagating in opposite directions). Hunter (1970, 1971) extended these results to an arbitrary nonstationary background flow, showing that cooling-dominated media are potentially more unstable than those in equilibrium, while heating provides stabilization. The most common applications of thermal instability to the ISM and star formation deal with the isobaric mode that was employed to explain the observed multiphase structure of the ISM \\citep{pikel'ner68,field..69,mckee.77,mckee90,wolfire....95}. Analysis of infinitesimal perturbations gives two characteristic length scales for the isobaric condensation mode: (1) a cooling scale $\\lambda_{\\rmn{p}}=c/\\omega_{\\rmn{p}}$ (where $c$ is the adiabatic sound speed and $\\omega_{\\rmn{p}}$ is the growth rate) and (2) a critical scale $\\lambda_{\\kappa}=c\\sqrt{t_{d}/\\omega_{\\rmn{p}}}$ (where $t_{d}$ is the characteristic thermal diffusion time). These two length scales define short-, intermediate-, and long-wavelength limits \\citep{meerson96,kovalenko.99}. In the short-wavelength limit, small isobaric perturbations are inhibited by heat conduction, so that $\\omega_{\\rmn{p}}<0$ for $\\lambda<\\lambda_{\\kappa}$. In the long-wavelength limit, the perturbations cannot grow isobarically because of the finite sound speed effects, and thus $\\omega_{\\rmn{p}}\\rightarrow0$ for $\\lambda/\\lambda_{\\rmn{p}}\\rightarrow\\infty$. This is only true if the gas is {\\em isochorically} stable \\citep{parker53,field65,shchekinov78}, otherwise the growth rate remains finite: $\\omega_{\\rmn{p}}\\rightarrow\\omega_{\\rmn{\\rho}}>0$ for $\\lambda/\\lambda_{\\rmn{p}}\\rightarrow\\infty$, but only large-scale temperature perturbations are growing, thus resulting in pressure variations and the formation of shock waves. The growth rates and characteristic scales depend on the heating and cooling properties of a given medium. Under the ISM conditions, if one assumes thermal equilibrium (i.e., an exact balance between cooling and heating), isochoric instability manifests itself only at relatively high temperatures, $T \\gsim 10^5$~K. However, in cooling-dominated regimes, it can develop at temperatures as low as $10^3$~K. \\begin{figure*} \\epsscale{2.0} \\plotone{fig1.ps} \\caption{Snapshots of the gas density field (perspective volume rendering): (a) First condensation at $t=0.07$~Myr, (b) thermal pancakes at $t=0.1$~Myr, (c) collapse and turbulization of cellular structure at $t=0.17$~Myr, (d) two-phase medium at $t=1.5$~Myr (5~pc box, $256^3$ grid points). The log density color coding is as follows: The dense blobs at the intersections of the filaments, $\\rho>10^{-22}$~g cm$^{-3}$, are light blue; the stable cold phase, $\\rho\\in[10^{-23}, 10^{-22}]$~g cm$^{-3}$, is blue; the unstable density regime, $\\rho\\in[10^{-23.7}, 10^{-23.0}]$~g cm$^{-3}$, is yellow to brown; and the low-density gas, including stable warm phase ($\\rho<10^{-23.7}$~g cm$^{-3}$), is a transparent red. The figure is also available as an mpeg animation in the electronic edition of the Astrophysical Journal. \\label{fig1}} \\end{figure*} A specific feature of TI in the ISM is a large ($\\gsim1$ dex) gap in gas densities between the two stable phases. The density range of interest in a galaxy formation context is even larger. This implies the importance of nonlinear effects in the dynamics of phase transitions. Nonlinearity brings into play nonequilibrium effects. Already weakly nonlinear development of condensations in an initially homogeneous gas in thermal equilibrium drives the system away from equilibrium. The mean pressure drops since $\\bar{\\rho^2} > \\bar{\\rho}^2$ and cooling overcomes heating globally \\citep{kritsuk85}. Later, on a timescale of $\\sim\\omega_{\\rmn{p}}^{-1}$, as condensations get denser and cooler, the isobaric condition $\\lambda\\ll\\lambda_{\\rmn{p}}$ becomes violated locally within them, so the system departs from pressure equilibrium. These effects are essential for the isobaric mode of TI in the ISM and star formation contexts. Therefore, analytical nonlinear solutions to ``isobarically'' reduced TI equations are insufficient to describe the radiative stage of the phase transition ($t\\sim\\omega_{\\rmn{p}}^{-1}$). During this strongly nonlinear stage {\\em large-scale} condensations form in such a way that gas moves almost inertially and its kinetic energy dominates thermal energy ($p\\ll \\rho v^2$). Accordingly, the gas velocities in these condensations are of the order of the sound speed in the unperturbed state. The situation here is directly analogous to the long-wave gravitational instability, so that results concerning the origin of cellular structure and \\citet{zel'dovich70} ``pancakes'' can be entirely carried over to the case of long-wave TI \\citep{meerson.87}. \\citet{sasorov88} gave an elegant proof that qualitatively the same result applies to {\\em small-scale} TIs; i.e., the onset of TI produces voids and highly flattened condensations along certain two-dimensional surfaces. These are also called ``thermal'' pancakes \\citep{meerson96}. The formation of filaments was simultaneously noticed in two-dimensional numerical simulations of TI in the solar transition region \\citep{dahlburg....87,karpen..88}. Ever since, thermal pancakes are being rediscovered both analytically and in numerical simulations (e.g., Lynden-Bell \\& Tout 2001). Thermal pancakes are transient. However, what happens next, before the evolution turns to a conductive relaxation stage \\citep{meerson96}, until recently has remained the ``terra incognito'' of TI theory. The problem of ``postradiative'' mechanical relaxation toward a static multiphase medium requires a solution for the full set of hydrodynamic equations that can only be obtained numerically. One-dimensional simulations pioneered by \\citet{goldsmith70} demonstrated that TI develops large motions in the ISM (see also Hennebelle \\& P{\\'e}rault 1999 for an example of how large motions can trigger TI). For some time, progress in this direction was precluded by numerical difficulties in modeling convergent cooling flows with high Mach numbers and high-density contrasts (e.g., V{\\'a}zquez-Semadeni, Gazol, \\& Scalo 2000). Recent multidimensional numerical simulations of the ISM evolution in disklike galaxies include effects of gravity, differential rotation, star formation and supernova feedback (de Avillez 2000; V{\\'a}zquez-Semadeni et al. 2000; Wada \\& Norman 2001; Wada 2001). However, it is hard to determine the role of TI in shaping the ISM structures found in these models, partly because the simulations still do not resolve length scales important for TI and partly because of the additional physical effects. The purpose of this Letter is to report on results of three-dimensional numerical simulations of classical TI that fill the gap in theory, exploring in detail the late radiative stage and postradiative relaxation toward a multiphase medium. Our major result is that formation of thermal pancakes induces turbulence in the ISM that serves as a nonlinear saturation mechanism for TI. As a consequence of a turbulent cascade, (1) information about initial perturbations is lost, including the imprints of heat conduction in the density power spectrum during the linear stage, and (2) turbulent diffusion becomes the dominant transport mechanism during the postradiative relaxation stage. \\begin{figure*} \\epsscale{1.8} \\vspace{-1.8cm} \\plottwo{fig2a.ps}{fig2b.ps} \\caption{Snapshots of phase diagrams (timing and labels correspond to those in Fig. \\ref{fig1}): (a) $t=0.07$~Myr, (c) $t=0.17$~Myr, (e) $t=0.5$~Myr, (d) $t=1.5$~Myr. The black dots show scatter plots of pressure vs. density. The ``dash'' at $P=4.55\\times10^{-10}$~dyn~cm$^{-2}$ in (a) shows the isobaric initial conditions. Background yellow-filled contours specify the part of the phase plane where isobaric mode is unstable; overlaid magenta contours are the regions of isochoric instability. The thick solid line shows thermal equilibrium curve. Density PDFs are plotted at the bottom of each panel (see scale to the right). \\label{fig2}} \\end{figure*} ", + "conclusions": "Our fiducial case was constructed to produce a stable two-phase medium because of its relevance to the Galactic ISM. We are interested in TI over a wide range of conditions as might be found in the ISM of high-redshift protogalaxies. We have simulated other cases with different parameter choices that do not produce stable two-phase media. However, we find they all develop turbulence in the nonlinear radiative stage of TI. Here we briefly discuss how the turbulence and asymptotic phase structure depend on initial conditions, deferring a more complete discussion to a future paper (A. Kritsuk \\& M. Norman 2002, in preparation). The level of induced turbulence is determined by the efficiency of conversion of the initial thermal gas energy into kinetic energy of turbulent flow by nonlinear development of TI: $E_k^{max}={\\cal{C}}(\\rho_0, T_0, \\varepsilon, Q, L)E_{th}(0)$ (see Fig. \\ref{fig3}). The conversion factor ${\\cal{C}}$ is a complex function of its variables. It varies from about 2\\% to $\\lsim1$\\% in our models. In general, higher $\\varepsilon$ and/or $T_0$ values provide higher conversion; a lower heating level ($Q<1$) supports TI and therefore works in the same direction. Larger boxes, as a rule, also produce more turbulence. The turbulence is induced on the initial cooling time and decays on a dynamical timescale, which is typically much longer. The turbulent Mach number at $t\\sim\\omega^{-1}_{\\rmn{p},0}$ depends on the mean temperature at this time. For nonequilibrium initial conditions ($Q\\ll1$ or $Q=0$), this is much lower than the initial temperature. In our fiducial case, ${\\cal{M}}_{rms}$ peaks at 8, dropping to 0.4 after 20 initial cooling times. For equilibrium initial conditions ($Q=1$), turbulent velocities remain subsonic (${\\cal{M}}_{rms}\\sim0.3$). But we would expect higher Mach numbers if the bistable range of pressure were wider than provided by our adopted cooling function. The initial gas density determines the number and mass fractions of thermal phases in the {\\em relaxed} state depending on its position relative to the valleys and hills on the thermal equilibrium curve. This is consequence of our choice of constant volume boundary conditions, which means that the mean density in the box remains constant. After the rapid cooling stage, our models with low initial densities $\\rho_0=1$-$5\\times10^{-25}$~g~cm$^{-3}$, high temperature $T_0=2\\times10^6$~K, and $Q\\in\\{0.3, 1\\}$ generate turbulence, evolve through a transient three-phase stage, and then relax to a single-phase low-pressure warm ISM. While turbulence is a generic feature of nonlinear saturation of TI, our simulations show that {\\em detailed} turbulent properties and the nature of emerging multiphase medium do depend sensitively on the Mach number and effective equation of state controlled by heating and cooling; this will be discussed elsewhere. Two identical simulations, except that cutoffs in initial power spectra were different ($k_{max}=8$ and 32 on a $128^3$ grid, $L=100$~pc, $Q=0$), demonstrated considerable structural differences in density distributions at the thermal pancake stage, $t_{tp}$, and surprisingly similar ``chaotic'' density structures and identical velocity power spectra at $\\sim6\\,t_{tp}$, when turbulent mixing covered the whole computational domain. This implies that the imprints of heat conduction in the density power spectrum during the linear stage could be erased later by the developing turbulent cascade. TI is certainly not the only potential source of turbulence in the ISM, but it cannot be ignored at least in those scenarios that actively employ TI to explain the origin and properties of observed objects. We suggest a paradigm shift concerning the role of thermal instability in the ISM and the nature of multiphase ISM. The idea of ``static'' two-phase ISM introduced in late 1960s (pressure-confined thermally stable dense clouds embedded in rarefied intercloud gas forming as a result of TI and subject to phase exchange due to cloud evaporation/condensation) must give way to the notion of a dynamic multiphase ISM, in which TI induces slowly decaying turbulence and in which turbulent diffusion regulates phase exchange processes. In this new emerging picture, the dense clouds are shapeless random aggregations of cold Lagrangian gas parcels; the clouds do not preserve their identity in real space on their sound crossing timescale until self-gravity tightens the fragments up into a self-gravitating cloud to form stars. Our results may suggests modifications to the scenario of a three-phase ISM \\citep{mckee.77,mckee90,heiles01a} that are yet to be understood." + }, + "0112/astro-ph0112092_arXiv.txt": { + "abstract": "{ A systolic algorithm rhythmically computes and passes data through a network of processors. We investigate the performance of systolic algorithms for implementing the gravitational $N$-body problem on distributed-memory computers. Systolic algorithms minimize memory requirements by distributing the particles between processors. We show that the performance of systolic routines can be greatly enhanced by the use of non-blocking communication, which allows particle coordinates to be communicated at the same time that force calculations are being carried out. The performance enhancement is particularly great when block sizes are small, i.e. when only a small fraction of the $N$ particles need their forces computed in each time step. Hyper-systolic algorithms reduce the communication complexity from $O(Np)$, with $p$ the processor number, to $O(N\\sqrt{p})$, at the expense of increased memory demands. We describe a hyper-systolic algorithm that will work with a block time step algorithm and analyze its performance. As an example of an application requiring large $N$, we use the systolic algorithm to carry out direct-summation simulations using $10^6$ particles of the Brownian motion of the supermassive black hole at the center of the Milky Way galaxy. We predict a 3D random velocity of $\\sim 0.4$ km s$^{-1}$ for the black hole. } \\begin{article} ", + "introduction": "\\label{sec:intro} Numerical algorithms for solving the gravitational $N$-body problem have evolved along one of two lines in recent years. Direct-summation codes compute the complete set of $N^2$ interparticle forces at each time step; these codes were designed for systems in which the finite-$N$ graininess of the potential is important, and are limited by their $O(N^2)$ scaling to moderate ($N\\lap 10^5$) particle numbers. The best-known examples are the {\\tt NBODY} series of codes introduced by Aarseth \\cite{Aarseth:99a}. These codes typically use high-order schemes for integration of particle trajectories and avoid the force singularities at small interparticle separations either by softening, or by regularization of the equations of motion \\cite{KS:65}. A second class of $N$-body algorithms replace the direct summation of forces from distant particles by an approximation scheme. Examples are tree codes \\cite{Barnes:89}, which reduce the number of direct force calculations by collecting particles in boxes, and algorithms which represent the large-scale potential via a truncated basis-set expansion (e.g. \\cite{Allen:90}) or on a grid (e.g. \\cite{Miller:68}). These algorithms have a milder, $O(N\\log N)$ scaling for the force calculations and can handle much larger particle numbers although at some expense in decreased accuracy \\cite{Spurzem:99}. The efficiency of both sorts of algorithm can be considerably increased by the use of individual time steps for advancing particle positions, since many astrophysically interesting systems exhibit a ``core-halo'' structure characterized by different regions with widely disparate force gradients. An extreme example of a core-halo system is a galaxy containing a central black hole \\cite{Merritt:99}. The efficiency of individual time steps compared with a global time step has rendered such schemes standard elements of direct-summation codes (e.g. \\cite{Aarseth:99b}). Here we focus on direct-summation algorithms as implemented on multi-processor, distributed-memory machines. Applications for such codes include simulation of globular star clusters, galactic nuclei, or systems of planetesimals orbiting a star. In each of these cases, values of $N$ exceeding $10^5$ would be desirable and it is natural to investigate parallel algorithms. There are two basic ways of implementing a parallel force computation for $O(N^2)$ problems. 1. {\\bf Replicated data algorithm.} All of the particle information is copied onto each processor at every time step. Computing node $i$, $1\\le i\\le p$, computes the forces exerted by the entire set of $N$ particles on the subset of $n_i=N/p$ particles assigned to it. 2. {\\bf Systolic algorithm.} At the start of each time step, each computing node contains only $N/p$ particles. The sub-arrays are shifted sequentially to the other nodes where the partial forces are computed and stored. After $p-1$ such shifts, all of the force pairs have been computed. \\noindent (The term ``systolic algorithm'' was coined by H. T. Kung \\cite{Kung:78,Kung:82} by analogy with blood circulation.) Both types of algorithm exhibit an $O(Np)$ scaling in communication complexity and an $O(N^2)$ scaling in number of computations. The advantage of a systolic algorithm is its more efficient use of memory: since each processor stores only a fraction $1/p$ of the particles, the memory requirements are minimized and a larger number of particles can be integrated. Other advantages of systolic algorithms include elimination of global broadcasting, modular expansibility, and simple and regular data flows \\cite{Kung:82}. The performance of a systolic algorithm suffers, however, whenever the number of particles on which forces are being computed is less than the number of computing nodes. This is often the case in core-halo systems since only a fraction of the particles are advanced during a typical time step. As an extreme example, consider the use of a systolic algorithm to compute the total force on a {\\it single} particle due to $N$ other particles. Only one processor is active at a given time and the total computation time is \\begin{equation} N\\tau_f + p(\\tau_l + \\tau_c) \\end{equation} where $\\tau_f$ is the time for one force calculation, $\\tau_l$ is the latency time required for two processors to establish a connection, and $\\tau_c$ is the interprocessor communication time. Thus the algorithm is essentially linear and no advantage is gained from having multiple processors. An efficient way to deal with the problem of small group sizes in systolic algorithms is via {\\it nonblocking communication}, a feature of MPI that allows communication to be put in the background so that the computing nodes can send/receive data and calculate at the same time \\cite{MPI:98}. In a nonblocking algorithm, the time per force loop for a single particle becomes \\begin{equation} {N\\tau_f\\over p} + p(\\tau_l + \\tau_c). \\end{equation} The second term is the waiting time for the last computing node to receive the particle after $p$ shifts. The first term is the time then required to compute the forces from the subset of $N/p$ particles associated with the last node. As long as the calculation time is not dominated by interprocessor communication, the speedup is roughly a factor of $p$ compared with the blocking algorithm. Here we discuss the performance of systolic algorithms as applied to systems with small group sizes, i.e. systems in which the number of particles whose positions are advanced during a typical time step is a small fraction of the total. Section 2 presents the block time step scheme and its implementation as a systolic algorithm. Section 3 discusses the factors which affect the algorithm's performance, and Section 4 presents the results of performance tests on multiprocessor machines of blocking and nonblocking algorithms. Section 5 presents a preliminary discussion of ``hyper-systolic'' algorithms with block time steps, which achieve an $O(N\\sqrt{p})$ communication complexity at the cost of increased memory requirements. Finally, Section 6 describes an application of our systolic algorithm to a problem requiring the use of a large $N$: the gravitational Brownian motion of a supermassive black hole at the center of a galaxy. ", + "conclusions": "We have introduced two variants of a systolic algorithm for parallelizing direct-summation $N$-body codes implementing individual block time step integrators: the first with blocking communication, and the second with non-blocking communication. Performance tests were carried out using $N$-body models similar to those commonly studied by dynamical astronomers, in which the gravitational forces vary widely between core and halo and for which the particle block sizes are typically very small. The nonblocking scheme was found to provide far better performance than the blocking scheme for such systems, providing a nearly ideal speedup for the force calculations. By engaging a sufficient number of computing nodes, particle numbers in excess of $10^6$ are now feasible for direct $N$-body simulations. For parallel machines with very large processor numbers, we describe the implementation of a hyper-systolic computing scheme which provides a communication scaling of $O(\\sqrt{p})$ at the expense of increased memory demands. The codes used to write this paper are available for download at:\\\\ \\verb+http://www.physics.rutgers.edu/~marchems/download.html+ \\begin{acknowledgment} This work was supported by NSF grant 00-71099 and by NASA grants NAG5-6037 and NAG5-9046 to DM. We thank Th.~Lippert, W.~Schroers, and K.~Schilling for their encouragement and help. We are grateful to the NASA Center for Computational Sciences at NASA-Goddard Space Flight Center, the National Computational Science Alliance, the Center for Advanced Information Processing at Rutgers University, the John von Neumman Institut f\\\"ur Computing in J\\\"ulich, and the H\\\"ochstleistungsrechenzentrum in Stuttgart for their generous allocations of computer time. \\end{acknowledgment}" + }, + "0112/astro-ph0112571_arXiv.txt": { + "abstract": "We have computed quasiequilibrium sequences of synchronously rotating compact binary star systems with constant rest masses. This computation has been carried out by using the numerical scheme which is different from the scheme based on the conformally flat assumption about the space. Stars are assumed to be polytropes with polytropic indices of $N=0.5$, $N=1.0$, and $N=1.5$. Since we have computed binary star sequences with a constant rest mass, they provide approximate evolutionary tracks of binary star systems. For relatively stiff equations of state ($N < 1.0$), there appear turning points along the quasiequilibrium sequences plotted in the angular momentum --- angular velocity plane. Consequently secular instability against exciting internal motion sets in at those points. Qualitatively, these results agree with those of Baumgarte {\\it et al.} who employed the conformally flat condition. We further discuss the effect of different equations of state and different strength of gravity by introducing two kinds of dimensionless quantities which represent the angular momentum and the angular velocity. Strength of gravity is renormalized in these quantities so that the quantities are transformed to values around unity. Therefore we can clearly see relations among quasiequilibrium sequences for a wide variety of strength of gravity and for different compressibility. ", + "introduction": "Binary neutron stars are very interesting objects. From the observational point of view, we will have a chance to get new eyes for the Universe by detecting gravitational waves in the first or second decade of this century. It is highly possible that the first signal may be that from compact binary stars, such as binary neutron stars, a black hole --- neutron star binary system, or binary black holes. On the other hand, theoretically, our understanding of evolution of compact binary stars is far from complete because it is considerably difficult to treat the ``2-body'' problem from a state with a wide separation to a merging stage consistently in the framework of general relativity. However, recent investigations have found a new approach to this problem. Since the time scale of the orbital change due to gravitational wave emission is rather long compared with the orbital period except for in the final few milliseconds of the coalescing stage, we can neglect gravitational wave emission for most stages of evolution. In other words, we can treat the system in ``quasiequilibrium'' (see e.g., \\cite{WMM96,BCSST98a,BGM99}). Following this idea, several groups have obtained quasiequilibrium sequences of binary neutron stars~\\cite{WMM96,BCSST98a,BGM99,UE00,UUE00}. Most of them adopted the assumption that the spatial part of the metric is conformally flat (the conformally flat condition: hereafter, CFC)~\\cite{WMM96,BCSST98a,BGM99,UE00}. For axisymmetric rotating polytropes results of the scheme with the CFC were compared with those obtained by the numerically exact code and found to be reasonably accurate~\\cite{CST96}. However, since there are no exact numerical solutions for binary configurations, one could not know the accuracy of the results obtained by the scheme with the CFC. Therefore, it is desirable to develop different schemes from that with the CFC and to compare results of different schemes for nonaxisymmetric configurations. As one of those alternatives, in the previous paper, we developed a new numerical scheme to obtain quasiequilibrium structures of nonaxisymmetric compact stars as well as the space time around those stars in general relativity and obtained quasiequilibrium sequences of synchronously rotating binary polytropes~\\cite{UUE00}. In that scheme, the Einstein equations are solved directly without assuming the CFC. The obtained results, however, could not be compared with those of \\cite{BCSST98a} because different polytropic relations were used. In this paper, we have used the same polytropic equation as that used in~\\cite{BCSST98a} and computed quasiequilibrium sequences of synchronously rotating polytropes with a constant rest mass. Therefore, we can directly compare our results with those with the CFC. In actual computations, we have improved our numerical scheme and succeeded in making our scheme more robust. ", + "conclusions": "In this paper we have solved quasiequilibrium sequences of synchronously rotating binary star systems in general relativity without assuming the CFC. We have constructed the constant rest mass sequences and shown that for stiff equations of state ($N < 1.0$), evolutionary curves have turning points so that synchronous rotation of the system breaks down at that point and that the internal motion will be excited. Our results can be compared with those of Baumgarte {\\it et al.} who employed the CFC~\\cite{BCSST98a}. Quantitatively, there are some differences between two results as seen from Figs.~\\ref{J-Omega N=1.0} and \\ref{J-Omega N=1.5}. These differences may come from different choices of the metric. However, it should be noted that qualitative features are very similar, i.e., the dependency on the polytropic index of the appearance of the turning points and so on. Therefore, although it is hard to give exact values of the angular velocity and/or the angular momentum at the turning points from quasiequilibrium approach, the occurrence of the instability could be correctly predicted. Nevertheless, since there are no exact solutions for the binary neutron star systems, we should keep in mind that there is a possibility that both of the two results might not represent the exact solutions. In Figs.~\\ref{dj-dm N=0.5} and \\ref{dj-dm N=1.0}, the nondimensional gravitational mass and the nondimensional angular momentum of each star are plotted against the nondimensional angular velocity. From these figures, it can be seen that turning points, i.e., the minima of each value, of two curves coincide. It implies that secular stability of binary systems can be found by investigating either the gravitational mass or the angular momentum. This is a nice feature that agrees the requirement between the change of the gravitational mass and that of the angular momentum as follows: \\begin{equation} dM = \\Omega dJ \\ , \\label{dj-dm} \\end{equation} where $dM$ and $dJ$ are the changes of the gravitational mass and the angular momentum of two configurations with the same rest mass, respectively. This relation can be reduced from the first law of thermodynamics, which is shown below, for the binary systems for which the rest mass, entropy, and vorticity of each fluid element are conserved~\\cite{FKM01}: \\begin{equation} dE = \\Omega dJ \\ , \\label{dj-de} \\end{equation} where $E$ means the half of the total energy of the system. It should be noted that these requirements can be checked if we can obtain highly accurate models. As seen from Figs.~\\ref{dj-dm N=0.5} and \\ref{dj-dm N=1.0}, changes of the gravitational mass and the angular momentum are three or four orders of magnitude smaller than the corresponding quantities. Unfortunately, since we cannot insist that our values have such high accuracy, we do not show our results here. As seen from Figs.~\\ref{J-Omega N=1.0}, \\ref{J-Omega N=1.5} and \\ref{J-Omega N=0.5}, the ranges of the values of $\\bar{M}_0 \\bar{\\Omega}$ and $\\bar{J}$ are considerably wide for the values of $N$. Even for the sequences with the same $N$, the values of $\\bar{M}_0 \\bar{\\Omega}$ and $\\bar{J}$ range widely. Thus it is not easy to understand the effects of the strength of gravity and/or the equation of state. In order to see the features of the evolutionary sequences at a glance, we introduce the following two nondimensional quantities, one of which can be considered to represent the angular velocity and the other of which corresponds to the angular momentum: \\begin{eqnarray} \\label{New J} \\hat{j} &\\equiv& \\frac{J}{J_0} \\ , \\\\ \\label{New omega} \\hat{\\omega} &\\equiv& \\frac{\\Omega}{\\Omega_0} \\ , \\end{eqnarray} where \\begin{eqnarray} \\label{normalization J} J_0 &\\equiv& \\frac{7}{5} G^{1/2} M^2 \\left(\\frac{M}{R}\\right)^{-1/2} = \\frac{7}{5} M \\frac{GM}{c^2} c \\left( \\frac{GM}{c^2R} \\right)^{-1/2} \\ , \\\\ \\label{normalization Omega} \\Omega_0 &\\equiv& \\frac{1}{2} G^{1/2} M^{-1} \\left(\\frac{M}{R}\\right)^{3/2} = \\frac{1}{2} c \\left(\\frac{GM}{c^2} \\right)^{-1} \\left(\\frac{GM}{c^2 R} \\right)^{3/2} \\ . \\end{eqnarray} Here, $R$ means the radius of the star on the major axis measured in the Schwarzschild-like coordinate. Our coordinate system in this paper is a kind of isotropic one and so $R$ is defined as follows: \\begin{eqnarray} R &=& r_{\\rm AB} \\left( 1 + \\frac{M}{2r_{\\rm AB}}\\right) ^ 2 \\ , \\\\ r_{\\rm AB} &=& \\frac{r_{\\rm B} - r_{\\rm A}}{2} \\ . \\end{eqnarray} The meaning of these quantities, ${\\hat j}$ and ${\\hat \\omega}$, can be roughly understood if we consider a system in Newtonian gravity which consists of two identical rigid spheres of uniform density in a contact phase. For such a system, $\\hat{j} = 1$ and $\\hat{\\omega} = 1$. Another property of these quantities can be seen from the definition of the normalization factors, Eqs.~(\\ref{normalization J}) and (\\ref{normalization Omega}). In these expressions, the differences originating from the different strength of gravity are ``renormalized\" by introducing the term related to the quantity $(M/R)_{\\infty}$. Thus we will call $\\hat{\\omega}$ a renormalized angular velocity and $\\hat{j}$ a renormalized angular momentum. In Figs.~\\ref{J-Omega New}, the renormalized angular velocity $\\hat{\\omega}$ is plotted against the renormalized angular momentum $\\hat{j}$ for several sequences of $N$ and $(M/R)_{\\infty}$. As seen from this figure, the values of $\\hat{\\omega}$ and $\\hat{j}$ for all evolutionary sequences with constant rest masses are scaled to values around unity. For the Newtonian sequences, the position of contact phases for smaller values of $N$ approaches $(1.0, 1.0)$ but never reaches that point because configurations are not rigid bodies and deformed from spheres by the tidal force from the companion star. Several characteristic features can be seen in this figure. First, if we compare the sequences with the same value of $(M/R)_{\\infty}$, sequences with stiffer equations of state locate at the upper-right region. This can be explained as follows. If we choose models which have the same values of the gravitational mass and the angular velocity, the radii are the same so that the inertial moment is larger for the stiffer polytropes. It implies that the angular momentum is larger for stiffer equations of state. Concerning the value of the renormalized angular velocity at the contact stage, the gravitational force is stronger for the stiffer polytropes because of the distribution of the matter inside the star. Thus larger angular velocity is required for configurations with stiffer equations of state. Second, if we compare the sequences with the same value of $N$, more relativistic sequences locate at the larger values of $\\hat{j}$. This is explained as follows. If we choose the models which have the same values of $N$, $M$ and $\\hat{\\omega}$, we obtain the following relation: \\begin{equation} \\hat{j} \\propto \\frac{< (r \\sin \\theta)^2 >}{R^2} \\ , \\end{equation} where $$ is the mass average of the quantity $F$ defined as follows: \\begin{equation} < F > \\equiv \\frac{\\int F dm}{\\int dm} \\ . \\end{equation} Here $dm$ is the mass element of the configuration. Since, in general, the change of the averaged quantity is smaller than the change of the quantity itself, the change of $\\hat{j}$ is affected mainly by the change of $R$ which decreases as $(M/R)_{\\infty}$ increases. Therefore the value of $\\hat{j}$ increases and the curves are shifted towards right in the plane. It should be noted that if we consider sequences with the same value of $(M/R)_{\\infty}$ but different values of $N$, differences due to different values of $N$ are amplified for configurations with the larger value of $(M/R)_{\\infty}$. Since real neutron stars cannot be approximated by a single polytropic relation all through the whole star, the evolutionary sequences cannot be approximated by the assumptions adopted in this paper, i.e., the assumption that $N$ and $K$ are conserved. Therefore, in order to get information about real evolutions, we need to construct evolutionary sequences with realistic equations. \\\\ We would like to thank Dr. K\\=oji Ury\\=u for his helpful discussions. FU is a Research Fellow of the Japan Society for the Promotion of Science (JSPS) and is grateful to JSPS for the financial support. This work was partially supported by the Grant-in-Aid for Scientific Research (C) of JSPS (12640255)." + }, + "0112/astro-ph0112217_arXiv.txt": { + "abstract": "The RossiXTE mission provided us with an unprecedentedly large database of X-ray observations of transient black-hole candidates. These systems are crucial for the understanding of the physical properties of mass accretion onto black holes. Here I review the results on a selected sample of systems and describe their behavior in a purely phenomenological way. From these results, we can derive a better classification of the spectral and timing characteristics of black hole candidates in terms of basic states. ", + "introduction": "Transient black hole candidates (TBHC) are the most important laboratories where we can study accretion onto black holes. The main reasons are two: first, there is a very limited number of bright persistent systems; second, transients go through a large range of mass accretion rate during their outbursts, therefore allowing us to study how the accretion properties change with accretion rate. In the case of persistent sources, they show rare state transitions (if any), and even when they do it is not clear what their dependence on accretion rate is (see e.g. \\cite{zhan1997}). Before the launch of the Rossi X-ray Timing Explorer (RXTE), a relatively small number of systems was known (see \\cite{tana1995,vdk1995} for a review). From these sources, a general ``canonical'' paradigm for the spectral/timing properties of TBHC had emerged (see \\cite{tana1995,vdk1995,miya1993}). Although there were notable exceptions, this paradigm was a good starting point for theoretical modeling. Four separate states were identified, which will be described here according to the behavior in the 2-20 keV band: \\begin{itemize} \\item Low/Hard State (LS): the energy spectrum can be described by a single power law with a photon index $\\Gamma\\sim$1.6. In addition, sometimes a weak disk-blackbody (DBB) component with kT$<$1 keV is observed, contributing less than a few percent to the detected flux. The Power Density Spectrum (PDS) is characterized by a strong band-limited noise with a break frequency below 1 Hz and a large fractional variability (30-50\\%). A low-frequency $<$1 Hz QPO is sometimes observed. \\item Intermediate State (IS): the energy spectrum can be decomposed in two components: a power law with $\\Gamma\\sim$2.5 and a clearly detectable DBB with kT$\\leq$1 keV. The PDS shows a band-limited noise with a break frequency higher than the LS (1-10 Hz) and a fractional variability of 5-20\\%. Sometimes a 1-10 Hz QPO is observed. \\item High/Soft State (HS): the energy spectrum is dominated by a DBB component with kT$\\sim$1 KeV, with the power-law component either below detection or extremely weak and steep ($\\Gamma\\sim$ 2--3). Very weak noise is detected in the PDS, in the form of a power-law with a few \\% of fractional variability. \\item Very High State (VHS): the energy spectrum is a combination of a DBB (kT$\\sim$1-2 keV) and a power law ($\\Gamma\\sim$ 2.5). The PDS can be of two types: either a band-limited noise similar to that of the IS, or a power-law, sometimes with a QPO. This state was observed only in two sources: the transient GS 1124-68 and the persistent source GX 339-4. \\end{itemize} The dependence of these states and their transitions on increasing accretion rate was determined mostly by the only transient system that had shown all four of them, GS 1124-68 \\cite{miya1994,ebis1994}: LS--IS--HS--VHS. As one can see from the description above, the IS is very similar to the VHS, both in energy and timing characteristics. What was believed to be different between them is the value of the accretion rate: the VHS was observed at very high accretion rates, while the IS was observed much later in the outburst of GS 1124-68, after a long period of HS, and therefore at lower accretion rate (see \\cite{bell1997}). Despite the exceptions, a transient was expected to have a fast-rise/exponential-decay light curve (reflecting the time history of accretion rate), lasting a few months, possibly with one or more re-flares, undergo a number of state transitions in the sequence outlined above following accretion rate changes, and return to quiescence after a period of a few weeks to months, until the next outburst. Once again, the prototypical source would be GS 1124-68. ", + "conclusions": "Since the launch of RXTE, thanks to the presence of the ASM and the operational flexibility of the mission, the available X-ray data on black hole candidates, especially transient systems, as increased substantially. The picture that emerges from the analysis of this large database is at the same time simpler and more complex than the paradigm that existed before. It is more complex because the phenomenology became quite complicated. The variety of QPOs observed in XTE J1550-564 and GRS 1915+105, and the extreme structured variability of the latter are perhaps the best examples. However, it also became somewhat simpler, as the number of basic spectral/timing states is now reduced to three, despite the tremendous diversity of some timing features, and the presence of a second parameter governing state transitions (although its physical nature needs to be addressed) might yield a direct measurement of a fundamental parameter of these systems. Among all sources, GRS 1915+105 shows more variability and state-transitions than all other sources together. it might be the way to solving basic problems of accretion, but it might also turn out to be an endless complication that leads away from the solutions. However, it is important that theoretical models address the general picture described above in addition to trying to describe in extreme detail the spectral distribution of a particular state in a particular source, especially since any spectral fit to low-resolution data involves by definition a strong a priori bias. \\small" + }, + "0112/astro-ph0112021_arXiv.txt": { + "abstract": "We are undertaking an extensive X-ray monitoring campaign of the two Crab-like pulsars in the Large Magellanic Clouds, PSR~B0540-69 and PSR J0537-6910. We present our current phase-connected timing analysis derived from a set of 50 pointed X-ray observations spanning several years. From our initial 1.2 yr monitoring program of the young 50 ms pulsar PSR~B0540-69, we find the first compelling evidence for a glitch in its rotation. This glitch is characterized by $\\Delta \\nu / \\nu = (1.90 \\pm 0.05) \\times 10^{-9}$ and $\\Delta \\dot \\nu / \\dot \\nu = (8.5 \\pm 0.5) \\times 10^{-5}$. Taking into account the glitch event, we derive a braking index of $n = 1.81 \\pm 0.07$, significantly lower than previous reported. For the 16 ms pulsar, PSR~J0537-6910, we recorded 6 large glitch events during a period of nearly 3 years, the highest rate of all known Crab-like systems. Despite the extreme timing activity, the long term spin-down of this pulsar continues to average $-1.9743 \\times 10^{-10}$ Hz/s. ", + "introduction": "A characteristic signature of young rotation-powered pulsars is the phenomena of ``glitches'', sudden discontinuous changes in their spin periods (e.g., see Lyne \\& Graham-Smith 1998). The physical causes of these glitches are not understood. Suggestions include sudden changes in the neutron star (NS) crust configuration (``starquakes''), abrupt reconfiguration of the magnetic field, or perhaps to the sudden unpinning of vortices in the superfluid neutrons in the inner part of the NS crust. For the latter, the amplitude of the glitch provides an estimate of the fractional part of the moment of inertia carried by superfluid neutrons (Lyne et al. 1996). The largest glitches have relative amplitudes ($\\Delta\\nu/\\nu$) of several parts per million, but the range of amplitudes covers many orders of magnitude. Often there is a partial recovery back toward the pre-glitch rotation rate on a time scale of $\\sim 100$ days, however, the spin-down rate may be permanently altered. Lyne (1995) noted that the amount of recovery in the rotation rate tends to be inversely proportional to the characteristic age of the pulsar. For some pulsars (e.g. Crab) the glitches are accompanied by a persistent increase in the spin-down rate with a relative amplitude of a few $\\times 10^{-4}$. This increase may be caused, for example, by changes in the alignment of the magnetic field because of starquakes (Allen \\& Horvath 1998; Link et al. 1998). In this paper we present preliminary results on the first detection of glitches from the two LMC Crab-like pulsars. ", + "conclusions": "" + }, + "0112/astro-ph0112351_arXiv.txt": { + "abstract": "{ We present the results of $B$, $V$, $R$ surface photometry of ESO~603-G21 -- a galaxy with a possible polar ring. The morphological and photometric features of this galaxy are discussed. The central round object of the galaxy is rather red and presents a nearly exponential surface brightness distribution. This central structure is surrounded by a blue warped ring or disk. The totality of the observed characteristics (optical and NIR colors, strong color gradients, HI and H$_2$ content, FIR luminosity and star-formation rate, rotation-curve shape, global mass-to-luminosity ratio, the agreement with the Tully-Fisher relation, etc.) shows that ESO~603-G21 is similar to late-type spiral galaxies. We suppose that morphological peculiarities and the possible existence of two large-scale kinematically-decoupled subsystems in ESO~603-G21 can be explained as being a result of dissipative merging of two spiral galaxies or as a consequence of a companion accretion onto a pre-existing spiral host. ", + "introduction": "The past several years have been very rich in observational studies of galaxy formation and evolution. One of the most interesting conclusions of these studies is the continuous assembly of galaxies (see Ellis 2001 for a recent review). Among the best local examples of delayed galaxy formation are the so-called multi-spin galaxies -- objects with more than one axis of rotation (Rubin 1994). Polar-ring galaxies (PRG) are probably the best known instance of multi-spin objects (see Whitmore et al. 1990, hereafter PRC, for definition and catalog of such objects ). PRG probably represent products of merger or external accretion phenomena (PRC, Reshetnikov \\& Sotnikova 1997). In this article we present the results of photometric observations of ESO~603-G21 -- a good PRG candidate according to Whitmore et al. (1990) (see Fig.~3f in Whitmore et al. 1990 and contour maps in our Fig.~2). This galaxy resembles an early-type galaxy with a well-developed bulge and an extended warped and edge-on disk/ring. A dust lane can be seen at the intersection of the bulge and the disk/ring. The spectroscopic data for this object indicate a complex scenario. The rotation curves for \\mbox{ESO~603-G21} show that the gas and stars in the disk/ring revolve around the minor axis (PRC, Arnaboldi et al. 1994). At P.A.~=~24$^{\\rm o}$ (minor axis), the spectra show no motion of the gas perpendicularly to the disk/ring. In contrast, the absorption line rotation curve indicates the existence of stellar motion along this axis (Arnaboldi et al. 1994). There stellar kinematics possibly indicate that the underlying stellar body is triaxial (Arnaboldi et al. \\cite{a1}, Arnaboldi et al. \\cite{a}). ", + "conclusions": "The global kinematical structure of ESO~603-G21 -- stellar rotation along two orthogonal position angles (Arnaboldi et al. \\cite{a1}) -- suggests that the object is a polar-ring galaxy. The host galaxy is probably an early-type galaxy with an exponential-like surface brightness distribution. The central galaxy is surrounded by a warped star-forming ring or disk. In general, ESO~603-G21 looks similar to other classic PRG (e.g. NGC~4650A). There are, nonetheless, several facts complicating such an interpretation. First, the central round component shows very low surface brightness which may indicate that the central object is not an early-type galaxy like in most classic PRG (PRC). Second, in the near-infrared region ($K$ passband) most of the stellar light comes from a bright nearly-exponential disk. Third, the central round object, clearly visible in the optical images (Fig.~2), is quite faint in the $K$ passband (Arnaboldi et al. \\cite{a}). The totality of the observed characteristics (optical and NIR colors, color gradients, HI and H$_2$ content, FIR luminosity and star-formation rate, rotation-curve shape -- Fig.~6 --, the agreement with the Tully-Fisher relation, etc.) suggests that ESO~603-G21 could be an unusual late-type spiral galaxy with a kinematically-decoupled extended \"bulge\". Therefore, it may be similar in some respects to NGC~4672 and NGC~4698, which are early-type disk galaxies with geometric and kinematical orthogonal decoupling between the bulge and disk (Bertola et al. \\cite{b}, Sarzi et al. \\cite{sa}), or to NGC~2748, which is a late-type spiral galaxy with possible ongoing accretion of a dwarf companion onto the central region of the galaxy (Hagen-Thorn et al. 1996). The bulge-like feature of ESO~603-G21 can be a \"true\" polar ring that is formed during almost perpendicular accretion of an early-type companion onto central region of a pre-existing disk galaxy. Another interesting interpretation of the observed ESO 603-G21 peculiarities is that the galaxy may be the result of a dissipative merger event (this scenario was proposed recently by Iodice et al. 2001 to explain the NGC~4650A puzzles). According to Bekki (1997, 1998), dissipative galaxy merging with a near polar orbit orientation can transform two late-type spirals into one PRG. In this scenario, a spiral galaxy falling from the polar axis of the target galaxy triggers the outwardly propogating density wave in the gaseous disk of the victim galaxy. Then, gaseous dissipation and star formation transform the victim disk into polar ring or disk. The central object is the intruding galaxy that has been turned into an early-type-like galaxy during the merging. Figure ~\\ref{draw} depicts the various internal substructures of ESO 603-G21 as revealed in this work (see item 3.2). Such complex, non-settled, fine structure of the galaxy supports our supposition about relatively late formation of the ''bulge'' due to external accretion or a merger. \\begin{figure} \\psfig{file=fig7.ps,width=8.7cm,clip=} \\caption[7]{Sketch of ESO 603-G21 as seen in careful visual inspection of the images in the various filters. North is up and East is on the left.} \\label{draw} \\end{figure} Interestingly enough, the companion of ESO 603-G21 is \\mbox{ESO 603-G20}, an edge-on disk-galaxy without any ``explicit\" evidence of interaction. The relative velocity between both objects is 65 kms$^{-1}$ (see NED and references therein); this suggests that both objects may form a bound system! There is, nonetheless, a third faint object between them, which seems to bear a very faint and narrow bridge to ESO~603-G21. The coordinates of the centroid (J2000) of this object are $\\alpha$ = 22$^{\\rm h}$ 51$^{\\rm m}$ 10.4$^{\\rm s}$ and $\\delta$ = -20$^{\\rm o}$ 14$'$ 59.5$\\arcsec$ within 1$\\arcsec$ of error, as calculated from the Digitized Sky Survey (DSS) image (see Fig.~8). Therefore, we have denoted this object Anon~J225110.4-201459.5. This is probably a low-surface brightness galaxy. On the basis of the DSS image we have found that the $B$-band total magnitude of the galaxy is $B_T=18.0\\pm0.5$. \\begin{figure} \\psfig{file=fig8.ps,width=8.7cm,clip=} \\caption[8]{In spite of the DSS resolution, a high-pass filtering has slightly enhanced the probable bridge between ESO 603-G21 (on the left) and Anon~J225110.4-201459.5 at the center (see text). ESO 603-G20 is on the right. North is to the top, east on the left.} \\label{dss} \\end{figure} It is essential to note that the disks of \\mbox{ESO 603-G21} (Fig.1) and \\mbox{ESO 603-G20} (Fig.~\\ref{dss}) are strongly warped. This feature, as well as the probable bridge, may be an indication of ongoing interaction in the \\mbox{ESO 603-G21}--\\mbox{ESO 603-G20}--Anon~J225110.4-201459.5 triple system (e.g. Reshetnikov \\& Combes 1999). So \\mbox{ESO 603-G21} is not an isolated object, but a member of a group of galaxies (like, for instance, NGC~4650A). Such dense spatial environment supports the idea that \\mbox{ESO 603-G21} may represent the result of a merging event. To test this scenario, detailed numerical simulations are needed." + }, + "0112/astro-ph0112398_arXiv.txt": { + "abstract": "{ In inhomogeneous optically thick synchrotron sources a substantial part of the electron population at low energies can be hidden by self-absorption and overpowered by high energy electrons in optically thin emission. These invisible electrons produce Faraday rotation and conversion, leaving their footprints in the linear and circular polarized radiation of the source. An important factor is also the magnetic field structure, which can be characterized in most cases by a global magnetic field and a turbulent component. We present the basic radiative transfer coefficients for polarized synchrotron radiation and apply them to the standard jet model for relativistic radio jets. The model can successfully explain the unusual circular and linear polarization of the Galactic Centre radio source Sgr A* and its sibling M81*. It also can account for the circular polarization found in jets of more luminous quasars and X-ray binaries. The high ratio of circular to linear polarization requires the presence of a significant fraction of hidden matter and low-energy electrons in these jets. The stable handedness of circular polarization requires stable global magnetic field components with non-vanishing magnetic flux along the jet, while the low degree of total polarization implies also a significant turbulent field. The most favoured magnetic field configuration is that of a helix, while a purely toroidal field is unable to produce significant circular polarization. If connected to the magnetosphere of the black hole, the circular polarization and the jet direction determine the magnetic poles of the system which is stable over long periods of time. This may also have implications for possible magnetic field configurations in accretion flows. ", + "introduction": "The detection of circular polarization (CP) in the continuum of radio sources is believed to be a powerful tool to test physical source models (Hodge \\& Aller \\cite{HA79}). But CP in extragalactic radio sources is extremely elusive (Roberts et al. \\cite{RC75}; Ryle \\& Brodie \\cite{RB81}; Weiler \\& de Pater \\cite{WdP83}) and is detected in only a few sources. A more recent ATCA-survey (Rayner, Norris \\& Sault \\cite{RNS2000}) for CP in radio-loud Quasars, BL Lacs and Radio Galaxies with improved sensitivity of $0.01\\%$, has shown a clear correlation of fractional CP with spectral index, in the sense that CP is stronger in flat and inverted spectrum sources. Circularly polarized radiation is therefore preferentially produced in self-absorbed radio cores. The fractional CP at 5 GHz is found to be between 0.05\\% and 0.5\\% in 11 out of 13 inverted spectrum sources at the ATCA spatial resolution of $2$ arcsec. At higher VLBA-resolution ($\\sim 0.5$mas) Homan \\& Wardle (\\cite{HW99}) report localized CP of 0.3\\%-1\\% in the jet-cores of 3C273, PKS 0528+134, and 3C279, which in a few cases may be as high as the local linear polarization. It is also found, that intraday variable sources are circularly polarized (Macquart et al. \\cite{MK00}), and that LP (linear polarization) and CP are both variable on timescales below 1 day. Recently CP was also found in X-ray binaries (Fender et al. \\cite{Fender2000}\\&\\cite{Fender2002}). While the handedness of CP is remarkably stable over several years (Komesaroff et al. \\cite{K84}; Homan \\& Wardle \\cite{HW99}; Fender et al. \\cite{Fender2002}) for individual sources, no preferred handedness of CP in general is found. An even more challenging situation than observed in radio-loud extragalactic jet sources presents itself in the centre of our galaxy. The compact radio source Sgr A$^*$ (see Melia \\& Falcke \\cite{MF01}), believed to be coincident with the dynamical centre of the Milky Way with a mass of $2.6\\,10^6 M_\\odot$ (Eckart \\& Genzel \\cite{EG96}; Ghez et al. \\cite{G98}) presumably in a single black hole, exhibits consistently larger circular than linear polarization in the range of $1.4$ to $15$ GHz (Bower et al. \\cite{BFB99}; Sault \\& Macquart \\cite{SM99}) with CP between $0.2$\\% and $1$\\%. LP is small and below the detection limits (Bower et al. \\cite{BB99}, \\cite{BW99}) with the exception of sub-mm measurements, which possibly shows LP at a level of $10\\%$ in the range $750-2000 \\mu$m (Aitken et al. \\cite{A2000}). The beam size of the sub-mm observations is $\\sim 10$ arcsec. The flux is dominated by extended dust emission or free-free emission and the synchrotron source is comparably weak at these wavelength. The inverted radio spectrum of Sgr A$^*$ ($S_\\nu \\propto \\nu^\\alpha,\\, \\alpha \\approx 0.3$) can be interpreted as either optically thin synchrotron emission (Beckert et al. \\cite{Beckert96}) or self-absorbed synchrotron emission from a jet-like outflow (Falcke et al. \\cite{FMB93}; Falcke \\& Markoff \\cite{FM00}). The idea of synchrotron emission by thermal electrons from Sgr A$^*$ was briefly considered by Reynolds \\& McKee (\\cite{RMcK80}) and revived for mildly relativistic electrons in the self-absorbed ADAF models for accretion in the galactic centre (Narayan et al. \\cite{N98}). The first ADAF models under-predicted the radio flux between $1$--$100$\\,GHz, which can be attributed to an outflow or jet. The upper limits for Sgr A$^*$ in the infrared require a sharp high energy cut-off for the electron distribution below $\\gamma_{\\mathrm{max}}$ of a few\\,$\\times 10^2$. Therefore thermal or quasi-monoenergetic electrons are responsible for the radio emission (Beckert \\& Duschl \\cite{BD97}), which distinguishes Sgr A$^*$ from high-luminosity, radio-loud AGNs. A close relative of Sgr A$^*$ is found in the centre of the normal spiral M81. The radio source M81$^*$ exhibits an elongated jet-like structure (Bietenholz et al. \\cite{BBR00}), has a similiar radio spectrum (Reuter \\& Lesch, \\cite{RL96}), a slightly larger luminosity, still below the AGN level, and has recently be found to be circularly polarized (Brunthaler et al. \\cite{Brun01}) without detectable LP. The fractional variability of CP is usually stronger than of LP, which in turn is stronger than for the total intensity. Together with the preserved handedness this poses servere constrains on possible scenarios for CP production and its variability (Komesaroff et al. \\cite{K84}). The suggested mechanisms are ({\\em a}) intrinsic cyclo-synchrotron emission from low-energy electrons or from electrons with small pitch angles seen close to the magnetic field direction (Legg \\& Westfold \\cite{LW68}), conversion from LP to CP as a propagation effect induced by ({\\em b}) low energy electrons inside the relativistic plasma (Hodge \\& Aller \\cite{HA77}) or ({\\em c}) by a magnetized cold plasma surrounding the synchrotron source. This requires either Faraday rotation (not possible in pure electron/positron jets) or changing (e.g., turbulent) B-field directions along the line of sight in the source. A further possiblity for CP production are ({\\em d}) inhomogeneous rotation measures in intervening cold plasma either close to the source or in our galaxy (Macquart \\& Melrose \\cite{MM2000}). The existence of these plasma screens can be infered from interstellar scattering believed to be the cause for intraday variability in some sources (Rickett et al. \\cite{R95}; Dennett-Thorpe \\& de Bruyn \\cite{DB00}; Macquart et al. \\cite{MK00}; Beckert et al. \\cite{BCore02}). This model predicts variable CP with a time averaged mean of $<$CP$> = 0$. In this paper we consider propagation effects like Faraday rotation and cyclic conversion of LP to CP and back (Pacholczyk \\cite{P73}) in turbulent, self-absorbed jets or outflows. First results were already published in Falcke et al. (\\cite{FB02}). We rederive some of the basic radiation transfer coefficients which, for example, could also be used for anisotropic particle distributions. The application of conversion to compact radio jets has been explored perviously by Jones (\\cite{J88}) using different techniques and without focusing on sources with large circular polarization and the role of globally ordered magnetic fields. Here we investigate the standard jet model with respect to the new polarization data placing some emphasis on the role of turbulence, the ratio of low- to high-energy particles, and the magnetic field confirguation. The paper is organized as follows: In Sec.\\,\\ref{emitrans} we review the basic production channels for CP. The outfow/jet model and the possible turbulence in the $B$-field is presented in Sec.\\,\\ref{outflowmod}. The consequences of Faraday rotation and conversion are discussed in Sec.\\,\\ref{dpolcon} followed by a detailed model of Sgr A$^*$. Polarization variability is the topic of Sec.\\,\\ref{polvar} and we close with a discussion of our results in Sec.\\,\\ref{discus}. ", + "conclusions": "\\label{discus} Recent observations of radio circular polarization in AGN, X-ray binaries, and the Galactic Centre black hole suggest that CP at the 0.3\\%-1\\%-level is common to many self-absorbed synchrotron sources. Faraday rotation and conversion in a magnetized and therefore bi-refringent plasma produce enhanced circular and reduced linear polarization. Both processes are sensitive to the presence of low-energy electrons and the orientation of the global magnetic field. The standard jet model for compact radio cores with a helical plus a turbulent magnetic field can well reproduce the circular and linear polarization spectrum of sources like Sgr A* and M81* with their high CP-to-LP ratio. The suppression of LP is achieved by the presence of a significant number of low-energy electrons in the source and an absence of an optically thin power-law extending to higher energies. The same model can also explain the typical level of circular polarization in blazars and the CP-to-LP ratio observed in blazars and X-ray binary jets. In this case a number of low-energy electrons is reduced with respect to the Sgr A* model and a power-law in the electron distribution exists. For Sgr A* the number of low-energy electrons producing conversion and depolarization needs to be significantly higher (by 2-3 orders of magnitude) than the number of radiating hot electrons. This means that a large fraction of the outflowing jet material is in the form of hidden matter shielded by self-absorption. This increases the estimates of the total jet power, which can be 5 orders of magnitude higher than the synchrotron luminosity. If one presumes that this power has to be provided by an accretion flow, the minimum accretion rates of $10^{-9..-8} M_{\\odot}$/yr, previously estimated from ``maximally-efficient'' jet models for Sgr A* (Falcke et al.~\\cite{FMB93}; Falcke \\& Biermann \\cite{FB99}) need to be raised to about $10^{-6} M_{\\odot}/$yr. This is quite consistent with recent estimates of Bondi-Hoyle accretion rates onto Sgr A* (Baganoff et al. \\cite{Baganoff2002}) and with suggestions for a coupled jet plus ADAF model (Yuan et al. \\cite{Yuan2002}), where the emission from the accretion disk is highly suppressed with respect to the jet. It is also interesting to note that to fit the CP with conversion one requires an asymmetry in the magnetic field components. This is naturally achieved by a helical magnetic field as is presumed to exist in jets. A symmetric configuration, e.g. a tightly wound helix or even a toroidal magnetic field structure would have difficulties to produce the observed level of CP. The stable handedness of CP over 20 years also implies a long-term stable component of the unidirectional field along the line-of-sight. This indicates that the polarity of the magnetic field (the ``magnetic north pole'') has remained constant over the last two decades. In view of the rather short accretion time scale in Sgr A$^*$ one could also speculate that this polarity is related to the accretion of a stable large-scale magnetic field which is accreted and expelled via the jet. The same can be said about blazars and X-ray binaries, where the stability found in GRS1915+105 by Fender et al. (\\cite{Fender2002}) is particularly interesting since the intrinsic accretion time scales in X-ray binaries are much shorter than those in supermassive black holes. Another important aspect of CP measurements is the question of the matter content of jets. We find that the constraints from CP of individual jet components for the jet power in blazars are not quite as severe as previously claimed and a statement in support of a pure electron/positron jet has to viewed with caution. For Sgr A* or M81* the situation may be different. If the depolarization is indeed intrinsic to the jet and not a surrounding medium (Agol \\cite{Agol2000}, Quataert \\& Gruzinov \\cite{Quataert2000}), one needs a high Faraday optical depth in the source, which can only be achieved by an excess of ``warm'' ($1\\la\\gamma\\la100$) electrons in an electron/proton plasma. While we have here assumed that all electrons are distributed in a single power-law, the actual situation may be quite different. For Sgr A* a power-law is actually not needed and we could obtain rather similar results with a two-temperature electron distribution, with temperatures corresponding to $\\gamma_{\\rm min}$ and $\\gamma_{\\rm max}$ respectively. This is not quite possible in blazars or bright X-ray binary jets, where extended electron power-laws are directly observed in the optically thin regime. It could well be that the radiative inefficiency of Sgr A* is due to the lack of effective shock acceleration that would increase the number of high-energy electrons with respect to the number of low-energy electrons (and in turn decrease the CP-to-LP ratio). The origin of these different electron distributions and their role for the radio-loudness of jet sources should be a very exciting question for further research. By improving our sensitivity and imaging all four Stokes parameters at multiple frequencies in the future, it will be possible to construct models of the entire emission and transfer processes in the source and determine the composition and energy spectrum of the relativistic plasma within jets. \\appendix" + }, + "0112/astro-ph0112167_arXiv.txt": { + "abstract": "Observations using the Australia Telescope Compact Array at a wavelength of 6~cm have uncovered the radio counterpart to the compact X-ray nebula surrounding the Vela pulsar. Two lobes were found oriented about the spin axis of the pulsar, starting at the edge of X-ray emission, they extend to three times the size. The northern lobe has a bright, defined edge and an integrated flux of 0.14 Jy, while the southern lobe of 0.12 Jy is more diffuse. ", + "introduction": "High resolution images from Chandra X-ray observations of the Vela pulsar have shown (Helfand, Gotthelf, \\& Halpern 2001; Pavlov et al. 2001) an X-ray pulsar wind nebula (PWN), which has been modelled in detail by Helfand et al. (2001) with an alternative interpretation by Radhakrishnan \\& Deshpande (2001). X-ray and synchrotron emissions are closely linked as highly charged particle flows drive both, so radio emission is a natural comparison. This targeted observation was optimised at 6~cm to match the Chandra image resolution and to improve on previous radio knowledge. We imaged a compact radio counterpart about the pulsar (figure 1). Previous radio studies (Bietenholz, Frail, \\& Hankins 1991; Frail et al. 1997a; Bock et al. 1998a; Bock, Turtle, \\& Green 1998b) have focussed on larger scales of filaments or wisps and the apparent connection to the Vela-X region of the supernova remnant. Reprocessed ATCA archive data showed diffuse extended emission in the region of the lobes and indicated a lack of emission in the X-ray PWN region, but it was not clear if this was due to sensitivity limits or a signature of radio emission. ", + "conclusions": "" + }, + "0112/astro-ph0112241_arXiv.txt": { + "abstract": "2MASSI J1315309$-$264951 is % an L3 dwarf with strong H$\\alpha$ emission discovered in the course of a color-selected survey for active galactic nuclei using % the Two-Micron All-Sky Survey (2MASS). The strength of its % H$\\alpha$ emission % decreased by about a factor of two between two epochs separated by 137 days. This is the first time that variable \\ha\\ emission has been reported in an L dwarf, and is probably the first observation % of an \\ha\\ flare in an L dwarf. The value of $\\log(L_{\\rm H\\alpha}/L_{\\rm bol}) > -4.17$ observed at the discovery epoch is larger than that of any other L dwarf but comparable to that of % 2MASSI J1237392+652615, the only reported T dwarf with \\ha\\ emission. The observed variability indicates that the \\ha\\ emission of 2MASSI J1315309$-$264951 is powered either by magnetic fields or by accretion in a binary system. Spectroscopic or narrow-band \\ha\\ monitoring of L and T dwarfs on timescales of hours to days would be the most useful step toward a better understanding of their \\ha\\ emission mechanism(s). % ", + "introduction": "\\label{INTRO} For many years the coolest, lowest-mass stars known were M dwarfs, but in recent years the L and T dwarfs have extended the stellar sequence to even lower temperatures and masses \\markcite{kir99}({Kirkpatrick} {et~al.} 1999). These three classes of low-mass stars and brown dwarfs exhibit distinctly different spectral features due to the decrease in photospheric temperature from M through L to T. In addition, M dwarfs often show H$\\alpha$ in emission. Down to the early M dwarfs, \\ha\\ activity correlates with rotation and thus decreases with age as stars lose angular momentum over time via stellar winds. Beyond objects of spectral type M8, however, it appears that \\ha\\ activity is stronger in more massive objects, even if they are older \\markcite{giz00}({Gizis} {et~al.} 2000). The frequency of \\ha\\ emission peaks around type M7 and declines for later-type L and T dwarfs, reaching zero at L5 \\markcite{giz00}({Gizis} {et~al.} 2000). However, contrary to this trend, \\markcite{bur00}{Burgasser} {et~al.} (2000) % reported the discovery of a T dwarf with strong H$\\alpha$ emission. Here I report the discovery of an L3 dwarf with similarly strong H$\\alpha$ emission. ", + "conclusions": "\\label{CONCL} 2MASSI\\,J1315309$-$264951 is an L3 dwarf with strong H$\\alpha$ emission which decreased in strength by about a factor of two between two epochs separated by 137 days, the first reported variable \\ha\\ emission in an L dwarf. The \\ha\\ emission in 2MASSI J1315309$-$264951 must be powered either by magnetic activity or by accretion in a binary system. Accreting binaries are rare, so that hypothesis is unlikely. The spectra presented here do not rule out a slow variation in \\ha\\ strength, but slow variations of the observed amplitude are rare among M dwarfs. Since flaring powered by reconnection of magnetic fields is common in M dwarfs and a radio flare has been detected in the L3.5 dwarf 2MASSW\\,J0036159+182110 \\markcite{ber02}({Berger} 2002), a flare is the logical explanation for the \\ha\\ variability in 2MASSI\\,J1315309$-$264951, The value of $\\log(L_{\\rm H\\alpha}/L_{\\rm bol}) > -4.17$ observed at the discovery epoch is larger than that of any other L dwarf but comparable to the value of $-4.3$ observed for 2MASSI J1237392+652615, the only reported T dwarf with \\ha\\ emission. However, both these values lie well below the average $\\log(L_{\\rm H\\alpha}/L_{\\rm bol})=-3.8$ observed in M dwarfs. Only two L dwarfs and one T dwarf are known to exhibit \\ha\\ emission of strength $\\log(L_{\\rm H\\alpha}/L_{\\rm bol}) > -5$. Thus perhaps two percent of L or T dwarfs exhibit \\ha\\ emission this strong at any given time. Given the small number statistics, this is consistent with the duty cycles observed for \\ha\\ flares ($\\sim$7\\%; \\markcite{giz00}{Gizis} {et~al.} 2000) and radio flares (2-10\\%; \\markcite{ber02}{Berger} 2002) among late-type dwarfs. Spectroscopic or narrow-band \\ha\\ monitoring of L and T dwarfs on timescales of hours to days is needed to determine if the frequency of strong \\ha\\ emission is governed by flaring, by the upper envelope of the magnetic field strength distribution, or by accretion in binary systems." + }, + "0112/astro-ph0112288_arXiv.txt": { + "abstract": "We measured the bar pattern speed, $\\om$, of the SB0 galaxy NGC 1023 using the Tremaine-Weinberg (1984) method with stellar-absorption slit spectroscopy. The morphology and kinematics of the \\hi\\ gas outside NGC 1023 suggest it suffered a tidal interaction, sometime in the past, with one of its dwarf companions. At present, however, the optical disc is relaxed. If the disc had been stabilized by a massive dark matter halo and formed its bar in the interaction, then the bar would have to be slow. We found $\\om = 5.0 \\pm 1.8$ \\kmsa, so that the bar ends near its co-rotation radius. It is therefore rotating rapidly and must have a maximum disc. ", + "introduction": "\\label{sec:introduction} Strong bars are seen in optical images of roughly 30\\% of all high surface brightness (HSB) disk galaxies (Sellwood \\& Wilkinson 1993) and this fraction rises to 50\\%-75\\% in the near IR (Knapen 1999; Knapen \\etal\\ 2000; Eskridge \\etal\\ 2000). Understanding the structure and dynamics of barred (SB) galaxies is, therefore, an issue of some importance. The principal dynamical quantity for SB galaxies is the rotation frequency/pattern speed of the bar, $\\om$. This is usually parametrized by the distance-independent ratio $\\vpd \\equiv \\lag/\\len$, where $\\len$ is the semi-major axis of the bar and $\\lag$ is the radius to the Lagrangian point, where the gravitational and centrifugal forces cancel out in the bar's rest frame. (The Lagrangian radius is therefore the generalization to strong bars of the corotation radius.) Contopoulos (1980) argued that a self-consistent bar is possible only if $\\vpd \\geq 1$. A bar is termed fast when $1.0 \\leq \\vpd \\ltsim 1.4$, while, for a larger value of $\\vpd$, a bar is said to be slow. A variety of methods have been used to attempt measurent of bar pattern speeds (see, for example, the review of Elmegreen 1996). Most measurements of $\\vpd$ rely on hydrodynamical simulations. These usually try to match the gas flow in the region of the bar, particularly at the shocks, which works because the location of the shocks depends on $\\vpd$, moving further ahead of the bar as $\\vpd$ increases. A bar needs to be fast for the shocks to remain curved with their concave sides towards the bar major axis, as is usually observed (van Albada \\& Sanders 1982; Athanassoula 1992). Detailed simulations of gas flows in individual galaxies also result in fast bars; examples include: NGC 1365 ($\\vpd = 1.3$, Lindblad \\etal\\ 1996), NGC 1300 ($\\vpd = 1.3$, Lindblad \\& Kristen 1996), and NGC 4123 ($\\vpd = 1.2$, Weiner \\etal\\ 2001). Hydrodynamical simulations can also recover $\\vpd$ by matching morphological features in \\hi; some examples are: NGC 7479 ($\\vpd =1.22$, Laine 1996), NGC 1073 ($\\vpd =1 - 1.2$, England \\etal\\ 1990), NGC 3992 ($\\vpd =1$, Hunter \\etal\\ 1989), and NGC 5850 $(\\vpd =1.35$, Aguerri \\etal\\ 2001). A direct method for measuring $\\om$, using a tracer population which satisfies continuity, was developed by Tremaine \\& Weinberg (1984). Since gas is subject to phase changes, it is not well-suited for this application. Old stellar populations in the absence of significant obscuration, on the other hand, are ideal for the Tremaine-Weinberg (TW) method. This has permitted application of the method to a small number of early type SB galaxies: NGC 936 ($\\vpd = 1.4 \\pm 0.3$, Kent 1987 and Merrifield \\& Kuijken 1995), NGC 4596 ($\\vpd = 1.15^{+0.38}_{-0.23}$, Gerssen \\etal \\ 1998) and NGC 7079 ($\\vpd = 0.9 \\pm 0.15$, Debattista \\& Williams 2001). The observational evidence, therefore, favors fast bars. The perturbation theory calculations of Weinberg (1985) predicted that a fast bar would be slowed down rapidly in the presence of a massive dark matter (DM) halo. Such slow-down has been seen in various simulations (Sellwood 1980; Little \\& Carlberg 1991; Hernquist \\& Weinberg 1992; Athanassoula 1996). The fully self-consistent, high resolution $N$-body simulations of Debattista \\& Sellwood (1998) also confirmed this prediction; however they showed that, for a maximum disc (here taken to mean a disc which dominates the rotation curve throughout the inner few disc scale-lengths, \\cf\\ van Albada \\& Sancisi 1986), a fast bar can survive for a large fraction of a Hubble time. Subsequently, Tremaine \\& Ostriker (1999) suggested that bars manage to remain fast not because discs are maximal but rather because the inner parts of DM halos are flattened and rapidly rotating. However, Debattista \\& Sellwood (2000) showed that rapid slow-down persists even then unless the halo angular momentum is very large relative to that of the disc. Thus they concluded that SB galaxies must be maximal, and argued that the same must be true for all high surface brightness disc galaxies. This conclusion rests on a small number of pattern speed measurements; in view of the fact that maximum discs are in conflict with the predictions of cold dark matter (CDM) cosmologies (\\eg\\ Navarro \\etal\\ 1997), enlarging the sample of measured pattern speeds is desireable. In this paper, we report observations of NGC 1023, for which we applied the TW method. The rest of this paper is organized as follows. The TW method is described briefly in \\S\\ref{sec:twmethod}. Then, in \\S\\ref{sec:properties} we give an overview of the previously known properties of NGC 1023. The photometric observations, reduction and results, including $\\len$, are presented in \\S\\ref{sec:photometry}, while \\S\\ref{sec:spectroscopy} presents the spectroscopic observations and results. We derive the rotation curve, corrected for the asymmetric drift, from which we obtain $\\lag$. With these results at hand, we then apply the TW method in \\S\\ref{sec:pattern_speed}. We present our conclusions in \\S\\ref{sec:conclusions}. ", + "conclusions": "\\label{sec:conclusions} We have found that the bar in NGC 1023 is fast, as are all bars which have been measured to date. Debattista \\& Sellwood (1998, 2000) showed that fast bars can persist only if the disc is maximal. Following Ostriker \\& Peebles (1973), it is sometimes thought that the unbarred (SA) galaxies are stabilized by massive DM halos. However, massive DM halos are not necessary for stabilizing discs; a rapidly rising rotation curve in the inner disc, such as when a massive bulge is present, is also able to inhibit bar formation (Toomre 1981; Sellwood \\& Evans 2001). Debattista \\& Sellwood (1998) argued that unbarred HSB galaxies must also be maximal for, if HSB disc galaxies form a continuum of DM halo masses spanning massive DM halo-stabilised SA galaxies to maximal SB galaxies, then slow bars must also be found in the intermediate range of halo masses. If we seek to avoid intermediate halo masses and slow bars by postulating (for whatever reason) a bimodal DM halo mass distribution for HSB galaxies, then we are left with the possibility that tidal interactions can still form bars, which would be slow (Noguchi 1987; Salo 1991; Miwa \\& Noguchi 1998). Thus Debattista \\& Sellwood (2000) concluded that the absence of slow bars requires that all HSB disc galaxies are maximal. However, it is possible that no such slow bars have been found because of an observational bias against SB systems with evidence of tidal interactions. We have chosen to study NGC 1023, in part, because it shows signs of a weak interaction in its past, without being at present significantly perturbed. The fast bar we found indicates that NGC 1023 has a maximal disc. If SA galaxies are stabilized by massive halos, we should find slow bars in that fraction of SB galaxies in which the bar formed through the interaction. While it is not possible to reach a general conclusion on the DM content of SA galaxies based on our measurement for a single galaxy, a large enough sample of similar SB galaxies with mild interactions in the past will be able to address this question. \\bigskip \\noindent {\\bf Acknowledgements.} \\noindent V.P.D. and J.A.L.A. acknowledge support by the Schweizerischer Nationalfonds through grant 20-56888.99. V.P.D. wishes to thank the Dipartimento di Astronomia dell'Universit\\`a di Padova for hospitality while preparing for the observations. E.M.C. acknowledges the Astronomisches Institut der Universit\\\"at Basel for the hospitality while this paper was in progress. We wish to thank the staff of the JKT telescope, particularly to the support astronomer J. C. Vega-Beltr\\'an and the staff of the TNG. We are indebted to R. Bender and R. Saglia for providing us with the FCQ package which we used for measuring the stellar kinematics. We thank the anonymous referee for suggestions that helped improve this paper. This research has made use of the Lyon-Meudon Extragalactic Database (LEDA) and of the NASA/IPAC Extragalactic Database (NED). \\bigskip \\noindent" + }, + "0112/astro-ph0112077_arXiv.txt": { + "abstract": "In this paper I argue that, far from necessarily hindering bar formation in disc galaxies, inner haloes may stimulate it. This constitutes a new instability mechanism by which bars can grow. To show this I use a number of $N$-body simulations whose initial conditions have identical discs and more or less concentrated haloes. They show that the bar that grows in the more halo-dominated environment is considerably stronger than the bar that grows in the more disc-dominated environment. This result is obtained from simulations with live haloes, i.e. composed of particles which respond to the disc and take part in the evolution. On the other hand, if the halo is rigid, it hinders or quenches bar formation, as expected. Comparison of two simulations which are identical in everything, except that the halo is live in the first one and rigid in the second one, leads me to suggest that the halo response can help the bar grow. Following the orbits of the stars in the halo, I find that a considerable fraction of the halo particles are in resonance with the bar. The halo may thus take angular momentum from the bar and stimulate its growth. I finally discuss whether and how the results of the $N$-body simulations can be applied to real galaxies. ", + "introduction": "It is by now well established that galactic discs can be bar unstable (e.g. Miller, Prendergast \\& Quirk 1970; Hohl 1971). In the quest for stability three main stabilising mechanisms have been proposed (see e.g. reviews by Athanassoula 1984, Sellwood \\& Wilkinson 1993 and references in either) : i) The disc could be immersed in a massive spheroid, e.g. a bulge and/or an inner halo (Ostriker \\& Peebles 1973). ii) The disc could be hot or have a hot center (Athanassoula 1983, Athanassoula \\& Sellwood 1986). iii) The galaxy could be sufficiently centrally concentrated to stop the initially linear wave from reaching the center (Toomre 1981; Sellwood \\& Evans 2001). Each one of these mechanisms has generated a lot of discussion, both regarding its efficiency and the way it can be applied to real galaxies. Here I will address the first of them, which was also historically the first to be introduced, and show that, contrary to what has been argued so far, inner haloes can, at least in some cases, enhance the bar. For this I will use a series of numerical simulations of disc-halo systems, described by Athanassoula \\& Misiriotis (2002, hereafter AM02) and Athanassoula (2002b in preparation, hereafter A02b). I describe the simulations and their results in section 2, and in section 3 I discuss the role of the halo. Finally in section 4 I give a general discussion and address the applicability of my results to real galaxies. ", + "conclusions": "\\label{sec:discuss} In the above I have compared three simulations starting off with identical discs, but different halo components. Any differences in their dynamical evolution should thus be attributed to the haloes. The strongest bar forms in the most halo dominated case, provided this is live, followed by the one in the disc-dominated case. In the simulation with the rigid halo there is only a very weak bar, or mild oval distortion, in the inner part. I thus reach the interesting conclusion that haloes can, at least in some cases, stimulate the bar instability and lead to stronger bars. This can be understood by a frequency analysis of the halo orbits, which reveals a large number of resonant orbits. Since these can exchange energy and angular momentum with stars at other resonances (Lynden-Bell \\& Kalnajs 1972) they can stimulate the bar instability, contrary to previous beliefs. The evolution of the galaxy leads to considerable concentration of the disc material in the central areas. Thus model MD starts off as disc dominated in the central parts, and, with time, the disc further enhances its superiority. Model MH starts off quite differently. Initially the halo is slightly more important than the disc within the radius at which the disc rotation curve is maximum, and considerably more so at larger radii, as witnessed from its circular velocity curve, shown in Fig.~\\ref{fig:inrotcur}. The central concentration of the disc increases considerably with time, so that, after the bar has grown, the disc dominates in the inner region. This may contribute an additional argument to the long standing debate of whether galactic discs are maximum or sub-maximum (e.g. Athanassoula, Bosma \\& Papaioannou 1987; Bosma 1999, 2000; Bottema 1993; Courteau \\& Rix 1999; Kranz, Slyz \\& Rix 2001; Sellwood 1999, Weiner, Sellwood \\& Williams 2001). Sackett (1997) and Bosma (2000) give a simple working definition to distinguish between maximum and sub-maximum discs, based on the value of $\\gamma = V_{d,max} / V_{tot}$, where $V_{d,max}$ is the circular velocity due to the disc component and $V_{tot}$ is the total velocity, both calculated at a radius equal to 2.2 disc scale lengths. According to Sackett (1997) this ratio has to be at least 0.75 for the disc to be considered maximum or maximal. In the simulations it is not easy to define a disc scale length after the bar has formed, so I will calculate $\\gamma$ at the radius at which the disc rotation curve is maximum, which is a well defined radius and is roughly equal to 2.2 disc scale lengths in the case of an exponential disc. Model MD starts off with $\\gamma > 0.75$, so that the disc starts maximum and stays so all through the simulation. In fact the value of $V_{d,max} / V_{tot}$ increases somewhat with time. Model MH has initially a value of $\\gamma$ around 0.68, i.e. close to the value of 0.63 advocated by Bottema (1993), and is therefore initially sub-maximum. This value, however, increases abruptly after the bar has formed, so that the disc can be considered maximum well before the time shown in Fig.~\\ref{fig:basic}, with a value of $\\gamma$ roughly equal to 0.86. Thus the formation of the bar leads the disc to evolve from sub-maximum to maximum, and hence strongly argues for maximum discs in disc galaxies with strong bars. This means that if we observe a strong bar in a disc galaxy the above simulations argue strongly and quantitatively that the underlying disc is maximum. The existence of gas should not alter this result. Indeed if the gas leads to a density distribution with a weak bar or no bar, then the above argument will be irrelevant, since it applies only to galaxies with strong observed bars. On the other hand, if the resulting bar is strong, then it should have rearranged the disc material sufficiently for the above argument to hold. We can reach similar results about the disc-to-halo mass ratio if we use the criterion of Athanassoula, Bosma \\& Papaioannou (1987), who examined what spiral perturbations can grow in a given disc/halo decomposition of an observed rotation curve. In a similar way I can calculate the $m$ component that will be strongest amplified via the swing amplification mechanism (Toomre 1981) in my simulations at or around the radius at which the disc rotation curve reaches its maximum. I find that, for model MD, it is the $m$ = 2 all through the simulation, as expected. The initial disc for model MH is certainly not maximum. I find that at $t$ = 0 higher $m$ components will be the most strongly amplified. The evolution, however, changes this, so that after the bar has grown it is the $m$ = 2 component that is the strongest amplified at or around the radius at which the disc rotation curve reaches its maximum. Both MD and MH models thus have, after the bar has grown, a disc which is intermediate between the ``no $m$ = 1'' and ``no $m$ = 2'' limits advocated by Athanassoula, Bosma \\& Papaioannou (1987) for real galaxies. Indeed these authors made a link between the structure present in a disc at a given time and the underlying halo mass at that time (not the initial halo mass) and thus their results are in good agreement with the above simulations, and many other similar ones (e.g. AM02 and A02b) . Apart from the halo, several other parameters can influence the formation of a bar. In particular let me stress the importance of the velocity dispersion of the disc particles, the effect of gas, and the effect of the velocity and mass distribution in the halo, as well as the existence of a gaseous companion. A complete description of all these effects is well beyond the scope of the present contribution and will be presented elsewhere. Let me just add a few preliminary words about the effect of the disc velocity dispersion. A sequence of MH-type galaxies shows that for larger initial velocity dispersion of the disc particles the bar is less strong and for sufficiently high values, becomes oval, or quasi-circular, in good agreement with what was found for 2D models (Athanassoula 1983). A sequence of MD-type models is more complicated. In these models the bar grows faster and becomes very long and strong. At that time, however, a strong buckling instability develops which leads to a considerable decrease of the bar amplitude. The final amplitude of the bar is a result of the competition between these two effects and this may be close. Only at sufficiently large values of the velocity dispersion can we be sure that the resulting oval will be very thick, as in the MH sequence. All the above are rather preliminary and will be discussed at length elsewhere. Finally the mass and velocity distribution of the halo component, together with the bar pattern speed and its time dependence, should influence how each of the resonance regions is populated and how responsive it is, and therefore influence its ability to exchange energy and angular momentum with the bar. Since very little is known on the composition of the halo, let alone about the distribution of the matter in it, it is very difficult to pursue this issue further. Nevertheless the arguments in section~3 lead to the prediction that at least some of the stars of the visible halo should be in resonance, in as much as they trace the relatively inner parts of the halo. Testing this would necessitate accurate information on the six phase space coordinates of a sufficiently large number of halo stars, as well as a sufficiently accurate description of the halo potential and the bar pattern speed. Our own Galaxy, which is barred, is the only place where advances with future astrometric satellites may make this possible, if we concentrate in areas which could have a high fraction of resonant stars." + }, + "0112/astro-ph0112307_arXiv.txt": { + "abstract": "We discuss the CM diagram of the galactic cluster NGC2420 to the light of current theoretical predictions. By relying on the most recent updating of the physical input, one finds too luminous theoretical He burning stars together with the evidence for a misfitting of the lower portion of the MS. Moreover one finds two well known overshooting signatures, as given by i) the large extension of the ``hook\" preceding the overall contraction gap, and ii) the scarcity of stars just at the end of the gap. We show that the overluminosity of He burning stars appears as a constant prediction of models based on updated physics, whereas alternative assumptions about the Equation of State can account for the MS fitting. Moreover, due to the scarse statistical significance of the observational sample, one finds that overshooting signatures can be present also in canonical (without overshooting) predictions. We conclude that, unfortunately, NGC2420 does not keep the promise to be of help in costraining the actual dimensions of convective cores in H burning MS stars, suggesting in the meantime that using clumping He burning stars as theoretical standard candle is at least a risky procedure. In this context the need for firmer constraints about the reddening of galactic clusters is shortly discussed. ", + "introduction": "The beautiful CM diagram presented in 1990 by Anthony-Twarog et al. for the intermediate age open cluster NGC2420 has been in the last ten years a favourite target for all the people concerned with the evolution of low to intermediate mass stars. The occasion for revisiting this cluster has been given to us by the recent paper by Pols et al. (1998, hereinafter P98), who presented new evolutionary tracks carefully discussing the fit of a selected (and well chosen) sample of galactic clusters. As discussed in that paper, NGC2420 seems to give a good chance to put firm constraints on the efficiency of the mechanism of core overshooting. Owing to the relevance of such an issue, we decided to go deep into the matter, hoping eventually to settle down such a long debated argument. However, as we will discuss in the following, the situation is far from being assessed. ", + "conclusions": "According to the previous sections, one finds that three out of the four theoretical misfitting of the cluster CM diagram can be accounted for within current evolutiony scenarios. On the contrary, no assumption appears able to reconcile the predicted luminosity of He burning stars with observation. To explore all the possibilities one may guess that a given amount of mass loss could account for this discrepancy. One generally assumes that mass loss occurs in the advanced phase of H shell burning, so that the internal structure of the He burning star is not affected by such an occurrence, which only decreases the amount of envelope surrounding the central He core. Under this assumption, the effect of mass loss on He burning models can be easily computed by simply decreasing the envelope of the constant-mass model. Numerical simulations shows that to reach the agreement between theory and observation for the clump luminosity one needs to decrease the He burning mass from the standard value of $\\approx$ 1.5 M$_{\\odot}$ to 1.0 M$_{\\odot}$. This very high amount of mass loss seems very unlikely to us. At present time the discrepancy between theory and observation in the luminosity of He burning stars with degenerate RGB progenitors it is not a surprising result; it appears as a constant prediction of models based on updated physics (see e.g. P98, Castellani et al. 2000), whereas similar suggestions for the overluminosity of theoretical models have been also derived from the pulsational properties of RR Lyrae (see e.g. Caputo et al. 2000). We conclude that using clumping He burning stars as theoretical standard candles, as we did (Castellani et al. 1999), is at least a risky procedure. We note that in the above quoted work, one could recognize a signature of the overluminosity in the need of assuming rather large reddenings for all the clusters. In this context, it follows that firm constraints about the reddening of galactic clusters will help in solving this problem." + }, + "0112/astro-ph0112131_arXiv.txt": { + "abstract": "We present a robust method to derive the duty cycle of QSO activity based on the empirical QSO luminosity function and on the present-day linear relation between the masses of supermassive black holes and those of their spheroidal host stellar systems. It is found that the duty cycle is substantially less than unity, with characteristic values in the range $3-6\\times 10^{-3}$. Finally, we tested the expectation that the QSO luminosity evolution and the star formation history should be roughly parallel, as a consequence of the above--mentioned relation between BH and galaxy masses. ", + "introduction": "The discovery of remarkable correlations between the masses of supermassive BHs hosted at the centers of galaxies and the global properties of the parent galaxies themselves~\\cite{8,9,12} leads to a natural link between the cosmological evolution of QSOs and the formation history of galaxies~\\cite{2}. The investigation of such interesting correlations looks promising not only to better understand how and when galaxies formed, but also to obtain information on the QSO population itself~\\cite{4}. Here we focus on two specific points raised by the general remarks above: 1) The use of the ``Magorrian relation'' to determine the QSO duty cycle at redshift $z=0$; 2) The expected relation between the cosmological evolution of the total luminosity emitted by star--forming galaxies and that of the total luminosity emitted by QSOs. As we will see, an interesting consequence of this last point is the possible existence of a physical process limiting gas accretion onto BHs at high redshifts. The {\\it observational} inputs of our analysis are the Magorrian relation~\\cite{12}, the galaxy mass--to--light ratio (from the Fundamental Plane)~\\cite{5,6}, the present-day luminosity function of spheroids~\\cite{14}, the present-day and the integrated QSO cosmological (light) evolution~\\cite{13}, and finally the star formation history~\\cite{11}. A possible alternative to the use of the mass-to-light ratio is the use of the Faber-Jackson~\\cite{7} relation coupled with the so-called $\\Mbh-\\sigma$ relation~\\cite{8,9}. The technical details will be given elsewhere~\\cite{4}. ", + "conclusions": "" + }, + "0112/astro-ph0112461_arXiv.txt": { + "abstract": "The BaseL Stellar Library (BaSeL) is a library of synthetic spectra which has already been used in various astrophysical applications (stellar clusters studies, characterization and choice of the COROT potential targets, eclipsing binaries, ...). This library could provide useful indications to 1) choose the best photometric system for the GAIA strategy by evaluating their expected performances and 2) apply systematically the BaSeL models for any sample of GAIA targets. In this context, we describe one of the future developments of the BaSeL interactive web site to probe the GAIA photometric data: an automatic determination of atmospheric parameters from observed colours. ", + "introduction": "The Basel Stellar Library (BaSeL) is a library of theoretical spectra corrected to provide synthetic colours consistent with empirical colour-temperature calibrations at all wavelengths from the near-UV to the far-IR (see Lejeune et al. 1997, 1998 for a complete description, and Westera et al. 1999 for the most recent version). These model spectra cover a large range of fundamental parameters (2000 $\\leq$ T$_{\\rm eff}$ $\\leq$ 50,000 K, $-$5 $\\leq$ [Fe/H] $\\leq$ 1 and $-$1.02 $\\leq$ log g $\\leq$ 5.5) and hence allow to investigate a large panel of multi-wavelength astrophysical questions, as briefly reviewed in the next section. Since they are based on synthetic spectra, they can in principle be used in many photometric systems taken either individually or collectively, and this is another major advantage of these models. The \"BaSeL interactive server\" is the web version of the BaSeL models ({\\tt http://www.astro.mat.uc.pt/BaSeL/}). This server is under development and the photometric systems presently available in interactive mode are: Geneva, Washington, Johnson-Cousins, Str\\\"omgren, HST-WFPC2, photographic RGU, and EROS. All details about this server will be given elsewhere. ", + "conclusions": "Covering a large spectral domain, extending from the UV to the far-infrared, the BaSeL models are adapted to perform simulations with the proposed GAIA photometric systems and should help to choose the most efficient one. We present and discuss a proposition to develop an automatic method, already used with success for COROT potential targets (Lastennet et al. 2001a), for a systematic determination of fundamental parameters from BaSeL synthetic multi-photometry. This new tool should be publicly available in 2002 on the following web site {\\tt http://www.astro.mat.uc.pt/BaSeL/}. \\small" + }, + "0112/astro-ph0112182_arXiv.txt": { + "abstract": "I review standard big bang nucleosynthesis and some versions of nonstandard BBN. The abundances of the primordial isotopes D, He-3, and Li-7 produced in standard BBN can be calculated as a function of the baryon density with an accuracy of about 10\\%. For He-4 the accuracy is better than 1\\%. The calculated abundances agree fairly well with observations, but the baryon density of the universe cannot be determined with high precision. Possibilities for nonstandard BBN include inhomogeneous and antimatter BBN and nonzero neutrino chemical potentials. ", + "introduction": "Big bang nucleosynthesis (BBN) is among the main observational evidence for big bang. The discovery of the cosmic microwave background (CMB) provided us with the temperature scale of the early universe, and allowed the calculation of the primordial nuclear abundances produced in the big bang. The four light isotopes, $\\D$, $\\EHe$, $\\UHe$, and $\\ZLi$ are mainly produced in the big bang, and the calculated abundances agree fairly well with astronomical observations. Standard big bang nucleosynthesis (SBBN) has a single free parameter, the baryon-to-photon ratio, \\begin{equation} \\eta \\equiv \\frac{n_b}{n_\\gamma} = 10^{-10}\\ldots10^{-9}, \\end{equation} which is related to the present baryonic contribution to the critical density $\\Omega_b$ via the Hubble constant $H_0 \\equiv h 100$~km/s/Mpc by \\begin{equation} \\et \\equiv 10^{10}\\eta = 274\\Omega_b h^2. \\end{equation} For decades, BBN has provided the best determination of the amount of baryonic matter in the universe. The agreement with observations is obtained in the range $\\et = 1.5\\ldots6$. Despite optimistic claims from time to time, BBN has not really progressed towards a much more precise determination of $\\eta$. Observers claim higher precision from determinations of primordial abundances of single isotopes, but disagree with each other or, within the context of SBBN, with primordial abundances of other isotopes. Difficult questions about systematic errors in observations and chemical evolution relating the present abundances to primordial abundances have prevented further progress. During the past year, a competing method for estimating the amount of baryonic matter has appeared. In the angular power spectrum of the anisotropy of CMB, the relative heights of the even and odd acoustic peaks are sensitive to the baryon-to-photon ratio. The first preliminary estimates from the Boomerang \\cite{Boom00} and Maxima-1 \\cite{Maxima00} experiments appeared to be in conflict with BBN, giving a higher baryon density, $\\Omega_b h^2 \\sim 0.03$, or $\\et \\sim 8$ \\cite{BooMax00}. The Boomerang collaboration has since revised their estimate downward, to $\\Omega_b h^2 = 0.022^{+0.004}_{-0.003}$ \\cite{Boom01}, which agrees with SBBN, but the Maxima-1 estimate has been revised upward to $\\Omega_b h^2 = 0.0325\\pm0.0125$ ($95\\%$ c.l.) \\cite{Maxima01}. With the coming satellite experiments CMB may surpass BBN as the method for estimating $\\eta$. BBN will then become a tool for understanding the astrophysics of chemical evolution, by telling us the primordial abundances. While SBBN is simple and natural, and is at present in reasonable agreement with observations, there is interest in studying nonstandard variants of BBN. For one thing, BBN is a sensitive probe of the physics of the early universe. If we change something about our assumptions regarding the conditions in the early universe, or the physics relevant at that time, we are likely to change the primordial abundances and ruin the agreement with observations. Thus for many things BBN provides the strongest constraint. On the other hand, from time to time there have been suggestions for disagreement between the estimated primordial abundances of the different isotopes, and/or other ways of estimating $\\eta$. If such disagreements persist, nonstandard BBN (NSBBN) may be the solution. I shall begin with a review of the physics of SBBN, and then discuss a few NSBBN scenarios. ", + "conclusions": "Standard BBN is compelling in its simplicity. While there is controversy among the observers and some apparent discrepancy between the estimated primordial abundances of the different isotopes and SBBN, these are probably not serious, and most likely represent difficulties in making observations and estimating primordial abundances based on observed ones. SBBN is thus not in trouble. Unfortunately, because of these difficulties, it is not able to pin down the baryon-to-photon ratio very precisely. It is somewhere in the range $\\et = 1.5$--$6$, or $\\Omega_b h^2 = 0.005$--$0.022$. The high redshift deuterium measurements point towards the upper end of this range, $\\et \\sim 5$--$6$, or $\\Omega_b h^2 \\sim 0.02$. The recent estimates from CMB anisotropy, $\\Omega_b h^2 = 0.022^{+0.004}_{-0.003}$ from Boom\\-erang \\cite{Boom01} and $\\Omega_b h^2 = 0.0325\\pm0.0125$ ($95\\%$ c.l.) from Maxima-1 \\cite{Maxima01} also favor this upper end of the range. We are eagerly waiting for more precise CMB measurements in the coming years. While standard BBN is in good shape, there is interest in studying nonstandard BBN: to assess the robustness of SBBN, to constrain possibilities for nonstandard physics or cosmology, and to be ready to provide relief if observational discrepancies turn out to be serious for SBBN. We discussed here four NSBBN scenarios: 1) electron neutrino degeneracy, 2) electron neutrino degeneracy combined with a speed-up of the expansion rate due to extra energy density, 3) inhomogeneous BBN, and 4) antimatter BBN. All these scenarios are able to relieve the tension between the $\\D$ and $\\UHe$ observations. The three last ones may also allow a larger baryon density than SBBN, but with some difficulty: 2) is constrained in that respect by other cosmological constraints and 3) and 4) cannot do much for the $\\ZLi$ constraint. I thank Elina Sihvola for permission to reproduce figures from \\cite{Sihvola01,Sihvola01a}." + }, + "0112/astro-ph0112527_arXiv.txt": { + "abstract": "The structure of dust spirals in the nuclei of the SAab-type Liner galaxies NGC 4450 and NGC 4736 is studied using archival HST PC images. The spirals are typically only several hundredths of a magnitude fainter than the neighboring disks, so unsharp mask techniques are used to highlight them. The ambient extinction is estimated to be less than 0.1 mag from the intensity decrements of the dust features and from the spiral surface filling factor, which is about constant for all radii and sizes. The nuclear dust spirals differ from main-disk spirals in several respects: the nuclear spirals have no associated star formation, they are very irregular with both trailing and leading components that often cross, they become darker as they approach the center, they completely fill the inner disks with a constant areal density, making the number of distinct spirals (the azimuthal wavenumber $m$) increase linearly with radius, and their number decreases with increasing arm width as a power law. Fourier transform power spectra of the spirals, taken in the azimuthal direction, show a power law behavior with a slope of $-5/3$ over the range of frequencies where the power stands above the pixel noise. This is the same slope as that found for the one-dimensional power spectra of HI emission in the Large Magellanic Cloud, and also the slope expected for a thin turbulent disk. All of these properties suggest that the dust spirals are a manifestation of acoustic turbulence in the inner gas disks of these galaxies. Such turbulence should dissipate orbital energy and transfer angular momentum outward, leading to a steady accretion of gas toward the nucleus. ", + "introduction": "Dust spirals in the inner kpc regions of galaxies reveal a source of compression that can affect the angular momentum distribution of the gas and possibly drive accretion to an AGN. Barred galaxies with an inner Lindblad resonance tend to have two long and symmetric dust spirals near the resonance that are a continuation of the leading-edge dust lanes in the bar (Athanassoula 1992). Many barred galaxies have ILR rings too (Buta \\& Crocker 1993; P\\'erez-Ram\\'irez et al. 2000; Knapen et al. 2000). Non-barred galaxies, galaxies without an ILR (e.g., late Hubble types), and regions of barred galaxies inside their ILRs can have more irregular dust spirals. Whether the nuclear spirals are regular or irregular, their presence in optical images suggests a density variation that is at least a factor of 2 and therefore likely to involve shocks. These shocks are oblique for azimuthal flows, so the gas will experience a torque when it enters a spiral and an opposite torque when it leaves. In cases where the gas moves faster than the spirals, the net torque is negative and the gas loses angular momentum and energy at the shock, causing it to spiral inward. If the spirals move faster than the gas, as might be the case outside a fast nuclear bar, then the gas gains angular momentum and may gradually move out. We are interested in the source of compression for the irregular nuclear dust spirals that appear in non-barred galaxies or inside the ILRs of some barred galaxies. We have proposed that some of these are caused by random sonic noise that amplifies weakly as it propagates toward the center (Elmegreen et al. 1998; Montenegro, Yuan, \\& Elmegreen 1999; Englmaier \\& Shlosman 2000). A signature of this process is an irregularity of structure, a wide range of pitch angles reflecting different times of origin and different wave propagation directions, a tendency to fill the volume available with a spiral separation comparable to the epicyclic radius for motions at the sound speed, and a general trend of increasing density with decreasing radius for each spiral. Because the velocities inside these spirals cannot yet be measured, there is no way to be sure they are driving inflow. Nevertheless, if any of the spirals come close to the nucleus, then such accretion would seem to be likely. Nuclear dust spirals and clouds have been studied previously using Hubble Space Telescope (HST) data. Van Dokkum \\& Franx (1995) found them on WFPC V-band images of early type galaxies. Malkan, Gorjian, \\& Tam (1998) did a large snapshot survey of active galaxies and classified the nuclear morphology including the dust. Elmegreen et al. (1998) noted dust spirals in WFPC2 images of the interacting galaxy NGC 2207 and proposed they might drive accretion to the nucleus. Regan \\& Mulchaey (1999) found nuclear dust spirals in 6 AGN galaxies using WFPC2 and NICMOS images, and Martini \\& Pogge (1999) found them with HST data in 20 Seyferts; both studies concluded that the spirals could drive accretion. Ferruit, Wilson \\& Mulchaey (2000) studied 12 early type Seyfert galaxies in more detail and noted the dust spirals too. Tomita et al. (2000) used HST archival images to study dust features in E and S0 galaxies, but did not comment on spirals specifically. Tran et al. (2001) studied nuclear dust in elliptical galaxies. Here we measure the extinction and structural properties of irregular dust spirals in the central kpc of two LINER galaxies, NGC 4450 and NGC 4736. Their Hubble types are about the same, SA(s)ab and RSA(r)ab (de Vaucouleurs et al. 1991), and their distances are taken to be 16.8 and 4.3 Mpc (Tully 1988). ", + "conclusions": "The nuclear spirals in NGC 4450 and NGC 4736 vaguely resemble the outer spirals, which are somewhat flocculent in each case, but the nuclear spirals do not continue smoothly from the outer spirals and there are important structural differences. The B-band image of NGC 4450 in the Carnegie Atlas of Galaxies (Sandage \\& Bedke 1994) shows two long dust spirals in the main disk, along with some flocculent structure; the stellar spiral arms are smooth. In contrast, the nuclear region has no stellar arms and at least 7 prominent dust arms, some with pitch angles as high as $45^\\circ$ and some crossing each other. The eastern side of the nuclear region shows more dust than the western side because of the galaxy's inclination. Some small dust feathers extend nearly radially from the center toward the south, reminiscent of jets. HST spectral observations by Ho et al. (2000) reveal double-peaked line profiles with high velocity wings, characteristic of accretion disk activity observed in other LINERs. NGC 4736 is an early-type galaxy with an outer ring and a circumnuclear starburst ring. Its main disk structure is flocculent and defined primarily by the dust. Sandage \\& Bedke's B-band print shows the inner disk structure as composed of many tightly-wrapped arms, but the central region is saturated in the reproduction. Waller et al. (2001) present UIT UV and ground-based R-band images of the central regions, including an unsharp-masked image showing the complicated flocculent structure of the main disk. They also show an HST FOC image of the main nuclear dust arms. Ground-based NIR observations by Mollenhoff, Matthias, \\& Gerhard (1995) suggested a weak stellar bar with a length of $20^{\\prime\\prime}$, which was also noted by Maoz et al. (1995) from an HST FOC image. The bar was observed in CO by Sakamoto et al. (1999) and Wong \\& Blitz (2000). In the HST image, the region corresponding to the bar shows up as an elongated disk with a position angle nearly perpendicular to the major axis of the galaxy. The structure inside the circumnuclear ring, which is really two tightly wrapped arms, consists of a dozen dust arms within a radius of 50 pc from the center, branching to dozens more dust spirals out to 200 pc. The nuclear dust spirals are not attached to the main inner disk dust spirals. The nuclear dust in NGC 4450 and NGC 4736 has several characteristics that differ from spiral arms and dust clouds in main galaxy disks. These are: \\begin{itemize} \\item The nuclear dust spirals shown here have no associated star formation. Other nuclear spirals in different galaxies have star formation (e.g., Coma D15 in Caldwell, Rose, \\& Dendy 1999), so the gaseous nature here is not universal. The lack of star formation suggests that the inner gas disks in NGC 4450 and NGC 4736 are not strongly self-gravitating. The same was true for NGC 2207 (Elmegreen et al. 1998) and for several other inner disks in the study by Martini \\& Pogge (1999). Our opacity estimate in Section 2.2 is also consistent with this. \\item The nuclear dust is in the form of spiral arms of various pitch angles, widths, and lengths. Some of the arms are trailing, a few are leading, and many cross each other. This pattern is generally more irregular than main disk flocculent arms (see atlas in Elmegreen 1981). Main disk flocculent arms are rarely leading. They generally do not cross each other; if they branch into spurs, then this branching is toward larger radii (Elmegreen 1980). They also have star formation that gives them a thicker, more patchy quality, rather than a filamentary quality. \\item The nuclear dust spirals in NGC 4450 and NGC 4736 have decreasing contrast with increasing radius. Ambient dust extinctions generally decrease with galactocentric distance because of the exponential distribution of gas column density. The flocculent galaxy NGC 5055 has such a decrease, for example, as measured by the extinctions of OB associations (Acarreta et al. 1996). Nuclear dust spirals are not just ambient extinctions, however. They are morphologically more similar to main disk spiral arms than diffuse cloud extinctions because they are organized and most likely formed by compressive processes in the presence of shear. From this point of view, nuclear dust spirals should be compared to main disk spirals, and then the radial decrease in nuclear spiral amplitude is unusual. Density wave spirals in non-barred galaxies tend to get stronger with increasing radius, out to at least the corotation zone (Elmegreen \\& Elmegreen 1984; Elmegreen et al. 1996). The unusual result that nuclear dust spirals get weaker with radius is presumably the result of crowding near the center for waves that move inward, as predicted for solutions to Bessel's wave equation (Elmegreen et al. 1998; Montenegro et al. 1999). This is a different dynamical situation than for main disk spiral arms, for which the curvature terms ($\\propto 1/kr$ for wavenumber $k$) in the wave equations can usually be ignored (e.g., Bertin et al. 1989). \\item The number density of dust features is about constant with radius, indicating that the inner disk is completely filled with structure. This is unlike the situation for main galaxy disks which often have a small number of arms (e.g., 2-5) that get further apart with radius. The nuclear spirals also have some indication of a hierarchical or fractal structure because of a non-integer slope of the size distribution, examined with unsharp masks. This size distribution is approximately a power law with a slope in the range from 0.5 to 0.8. This power law is reminiscent of other properties of interstellar clouds formed by turbulence, but complicated in this case by the effects of shear, which make spirals rather than clumps, and by the Coriolis force, which resists turbulent motions on large scales. The dust features are also very weak, and the smaller clouds, as well as those further from the center, are difficult to see above the pixel noise. Fractal structure in the dust of another galaxy was also found by Keel \\& White (2001) using a background illumination source. Fractal structure in galactic clouds is well known (e.g., Falgarone, Phillips, \\& Walker 1991). \\item Fourier transform power spectra in the azimuthal direction show the characteristic signature of turbulence compression, which is a power law slope of $-5/3$ for one-dimensional structures that are larger than the line-of-sight thickness of the galaxy disk. This makes the nuclear dust features studied here resemble the HI clouds in the LMC, with the important difference that the nuclear clouds are spiral filaments, presumably affected by shear, and the LMC clouds are globular and shell-like in a low shear environment. \\end{itemize} There is no direct evidence in our observations for accretion driven by the dust spirals. Radial velocities will have to be measured to determine this. However, the increase in dust opacity with decreasing radius for the main spirals is consistent with the amplification that is expected for inward motions. In that case, the spirals could drive nuclear accretion." + }, + "0112/astro-ph0112019_arXiv.txt": { + "abstract": "We report on a long-term monitoring campaign of \\axp, the 12-s anomalous X-ray pulsar and magnetar candidate at the center of the supernova remnant \\kes73. We have obtained approximately monthly observations of the pulsar with the {\\it Rossi X-ray Timing Explorer} (\\xte) spanning over two years, during which time \\axp\\ is found to be rotating with sufficient stability to derive a phase-connected timing solution. A linear ephemeris is consistent with observations of the pulse period made over the last 15~yrs with the \\ginga, \\asca, \\xte, \\& \\sax\\ observatories. Phase residuals suggest the presence of ``timing noise'', as is typically observed from young radio pulsars. These results confirm a rapid, constant spin-down for the pulsar, which continues to maintain a steady flux; this is inconsistent with most accretion scenarios. ", + "introduction": "1E~1841$-$045 is perhaps the best candidate for a magnetar -- an isolated neutron star (NS) with an extreme magnetic field whose energy loss is dominated by magnetic field decay (Vasisht \\& Gotthelf 1997; see Duncan \\& Thompson 1992 for theory of magnetars). This interpretation is based primarily on initial studies which showed that the pulsar is slowing down rapidly ($\\dot P = 4.1 \\times 10^{-11}$ s/s), with an inferred equivalent magnetic dipole field of $B_p \\sim 7.0 \\times 10^{14}$ G, nearly twenty times the quantum critical field, while producing steady emission at a rate far in excess of its rotational kinetic energy loss (Vasisht \\& Gotthelf 1997). Further evidence for an isolated system is the absence of Doppler shifts, detectable companion or accretion disk, and a spectrum which differs greatly from those of known accretion-powered binary systems. Here we report on a two-year monitoring campaign of \\axp\\ made with \\xte. We have obtained a spin ephemeris derived from a phase-connected timing solution which is found to be consistent with nearly two decades of observations of the pulsar. We are able to characterize rotational stability of \\axp, including a search for timing noise, pulsed flux variability, and frequency glitches. These results provide important observational constraints on both the magnetar and any accretion model for \\axp. ", + "conclusions": "Observationally \\axp\\ is characterized as an Anomalous X-ray Pulsar (AXP; see Mereghetti 2001 for a review). This small group ($\\sim 6$ members) of seemingly isolated pulsars are observed exclusively in the X-ray energy band\\footnote{There is recent evidence for IR emission from two AXPs, see Hulleman et al 2000, 2001.}, are underluminous for an accreting X-ray binary system, and are, except for the characteristic burst activity, otherwise similar observationally to the SGRs. This is the fourth AXP, after 1E~2259+586, 1E~0142+615, and RXS J1708$-$4009, for which long-term phase-connected timing is possible, characterizing these objects as rotators of great stability (see Gavriil \\& Kaspi 2001). In contrast, the AXP 1E~1048.1$-$5937 is much less stable, as a phase-connected timing solution cannot be maintained for more than a few months (Kaspi et al. 2001). The measured strength of the torque noise for \\axp\\ is comparable to that found for 1E 1048.1$-$5937, during its intervals of relative stability, but substantially weaker than the torque noise measured for most accreting pulsars (Bildsten et al. 1997). Still, it is a factor of two quieter than even the most exceptionally quiet accreting pulsars (e.g., 4U 1626$-$67, Chakrabarty et al. 1997). The leading theory for the nature of AXPs is the magnetar model as first proposed by Thompson \\& Duncan (1996). In this model, in the absence of soft-gamma-repeater-like outbursts, one expects generally smooth spin-down. The lack of deviations from a simple spin-down model found for \\axp\\ is consistent with the magnetar model. An alternative model proposed for AXPs is that they are accreting from a disk of material that formed shortly after the supernova explosion that gave birth to the neutron star. Severe constraints have been placed on the plausibility of this scenario by optical/IR observations of AXPs (Hulleman et al. 2000). Both the constant, long-term spin-down of \\axp, as well as the fact that phase-connected timing is possible for this source, provide evidence against accretion scenarios, as spin-up episodes, as well as considerably more torque noise, are in general to be expected." + }, + "0112/astro-ph0112196_arXiv.txt": { + "abstract": "{ Within a project to investigate the properties of \\LB\\ stars, we report on their abundance pattern. High resolution spectra have been obtained for a total of twelve candidate \\LB\\ stars, four of them being contained in spectroscopic binary systems, and detailed abundance analyses have been performed. All program stars show a characteristic \\LB\\ abundance pattern (deficient heavy elements and solar abundant light elements) and an enhanced abundance of Na. This work raises the fraction of \\LB\\ stars with known abundances to 50~\\%. The resulting abundances complemented by literature data are used to construct a ``mean \\LB\\ abundance pattern'', which exhibits, apart from general underabundances of heavy elements ($\\approx-1$~dex) and solar abundances of C, N, O, Na and S, a star-to-star scatter which is up to twice as large as for a comparable sample of normal stars. ", + "introduction": "\\LB\\ stars are defined as Population\\,{\\sc i} A- to F-type stars, which are metal-poor but exhibit nearly solar element abundances for C, N, O and S \\citep[e.g.][]{Hauc:83,Abt:84b,Gray:97,Paun:00}. A historical review of studies dealing with these stars, as well as extensive work on several aspects regarding the \\LB\\ group, like spectral classification, photometry, ultraviolet and infrared fluxes, space motions, binarity and theoretical considerations, can be found in \\citet{Paun:00}. From the definition above it is obvious that the chemical composition of the atmospheres of \\LB\\ stars represents the most important property for the characterization of this group of stars. At the beginning of this project, few abundance determinations for \\LB\\ stars existed in the literature (see Sect.~\\ref{pattern}). It was not possible to make a firm conclusion about ``the abundance pattern'' of the \\LB\\ stars, in particular because only a part of the parameter space occupied by current lists of \\LB\\ candidates was covered by abundance analyses. A detailed knowledge of the common abundance pattern, if existent, and of the variation of abundances of individual elements with stellar parameters is the key to understanding the astrophysical processes generating the properties of the \\LB\\ group. Therefore the determination of abundances for as many elements as possible for preferably all members of this group (where the membership assignment is based on classification spectroscopy) is regarded as a main goal of this investigation. Together with all abundance analyses found in the literature, the proportion of \\LB\\ stars with known abundances, based on the list of \\LB\\ stars given by \\citet{Paun:00}, amounts now to about 50\\%, which allows to draw conclusions on the abundance pattern of the \\LB\\ stars as a group. A statistical investigation of the \\LB\\ parameters and a comparison with other groups of stars and the interstellar medium will be the topic of a subsequent paper. \\begin{table*} \\caption{Observations used in this work. The last three columns give the observed wavelength range ($\\lambda$), the resolving power (R), and the mean continuum signal to noise ratio determined from the reduced spectra.} \\label{obs_tab} \\begin{tabular}{llclccl} \\hline\\hline HD & Observatory / Telescope & Date & Observer & $\\lambda$ [\\AA] & R & S/N \\\\ \\hline\\hline 74873, 84948, 101108, 106223, & Asiago (Italy) / 1.8 m & 3--1995 & U.\\,Heiter, & 4000$-$5700 & 20000 & 150, \\\\ 110411 & & & E.\\,Paunzen & & & 250 \\\\ 84948, 171948 & & 2--1997 & & 4500$-$7200 & & 200 \\\\ \\hline 107233, 142703, & OPD/LNA (Brazil) / 1.6 m & 6--1995 & E. Paunzen & 3900$-$4900 & 22000 & 150, \\\\ 168740, 170680 & & & & & & 200, 230 \\\\ \\hline 106223 & KPNO (Arizona) / 1.5 m & 4--1998 & M. Weber & 4240$-$4570 & 24000 & 200 \\\\ \\hline \\hline 74873, 101108, 170680 & DSO (N. Carolina) / 0.8 m & 1995 & R.O.\\,Gray & 4000$-$4450 & 2200 & 300, 150, 400 \\\\ \\hline 106223 & OHP (France) / 1.93 m & 1995 & E.Paunzen, & 4000$-$4450 & 1200 & 300 \\\\ \\cline{1-3}\\cline{5-7} 110411 & Asiago (Italy) / 1.8 m & 1995 & \\ U.Heiter & 4000$-$4450 & 1100 & 350 \\\\ \\hline 107233, 142703 & OPD/LNA (Brazil) / 1.6 m & 1995 & E.Paunzen & 4000$-$4450 & 2200 & 200, 350 \\\\ \\hline\\hline \\end{tabular} \\end{table*} ", + "conclusions": "We have presented new abundances of up to 15 elements for twelve \\LB\\ stars, which include five Y abundances. This element has up to now been studied only for two stars: HD\\,84123 \\citep{Heit:98} and 29\\,Cyg \\citep{Adel:99}. In these stars, the Y abundances are enhanced with respect to that of most other heavy elements. The same is found for all stars presented here. The other heavy elments show underabundances of varying degree, up to about $-$2~dex, and the light elements C, O and Na appear moderately over- or underabundant, except for HD\\,171948A/B. In this SB2 system [C] and [O] are less than $-$0.5~dex, but this is still high compared to the Fe abundance of $-$1.6~dex. The results of this work consolidate the membership of all candidate stars to the group of \\LB\\ stars, nearly doubling the amount of spectroscopically investigated members of this group of peculiar stars. In addition, we have collected abundances for 34 \\LB\\ stars from the literature. They were used to construct a ``mean \\LB\\ abundance pattern'', which can be described as follows: \\begin{itemize} \\item The iron peak elements from Sc to Fe as well as Mg, Si, Ca, Sr and Ba are depleted by about $-$1~dex relative to the solar chemical composition. The mean abundance of Zn is similarly low, although its condensation temperature is similar to that of S and should prevent the depletion of this element by chemical separation within an accretion scenario. \\item Al is slightly more depleted ($-$1.5~dex) and Ni, Y and Zr are slightly less depleted. \\item The abundances of the light elements C, N and O as well as S lie around the solar value, which is not surprising because this is part of the definition of the \\LB\\ group. The mean abundance of Na is also solar, but the star-to-star scatter is much larger for this element ($\\pm$1~dex). \\item The star-to-star scatter is twice as large as for a comparable sample of normal stars for all elements except Ba, which means that for more than half of the heavy elements at least one star is included in the sample, for which the abundance of one or more of these elements is solar. \\end{itemize} In contrast to the results of \\citet{Paun:99a}, C is on average {\\em more} abundant than O when the abundances are determined from optical lines. The enhanced abundance of Na has been already noted by \\citet{Stue:93}, who derived a mean abundance of +0.6~dex for his sample of stars, and by \\citet{Paun:99b}. This description suggests the existence of a separate chemically peculiar group of ``\\LB\\ stars'' with a characteristic abundance pattern. On the other hand, the scatter of abundances for each element indicates that the ``\\LB\\ group'' is rather inhomogeneous. Furthermore, the comparison to normal stars is difficult because the sample of ``normal'' main sequence stars with known abundances and parameters similar to that of the \\LB\\ stars is rather limited." + }, + "0112/astro-ph0112475_arXiv.txt": { + "abstract": "A measurement of the lateral distribution of Cherenkov light in Extensive Air Showers (EAS) at $E_\\circ$ = 10 $\\div$ 100 TeV and a study of the compatibility of the photon number spectra with the expectations from the direct measurements of $p$ and $\\alpha$ spectra and the Corsika-QGSJET propagation code in the atmosphere have been performed at the National Gran Sasso Laboratories by the EAS-TOP and MACRO arrays. The telescope array of EAS-TOP has been used as the Cherenkov light detector. The muon tracking system of MACRO in the deep underground Gran Sasso Laboratories ($E_{\\mu}$ $>$ 1.3 TeV) served as the EAS detector, including core localization and arrival direction. ", + "introduction": "The EAS-TOP and MACRO arrays at the Gran Sasso Laboratories offer a unique opportunity of measuring the lateral distribution of Cherenkov light in the 10 $\\div$ 100 TeV energy range by associating the Cherenkov light collected by the EAS-TOP telescopes with the TeV muon reconstruction, and consequently the EAS core geometry, through the MACRO array. In this paper we report on the measurements of the Cherenkov light lateral distribution compared with the results of simulations based on the CORSIKA-QGSJET code providing an experimental validation of the code itself. Moreover the technique allows a study of the primary composition and a comparison with the direct existing measurements in a overlapping region. Due to the shower selection through the high energy muon ($E_{\\mu} >$ 1.3 TeV, i.e. primary energy $E_{o} >$ 1.3 TeV/nucleon), in the energy range 10 TeV $<$ $E_{o}$ $<$ 40 TeV (10 TeV being the Cherenkov telescopes' threshold energy) the selected primaries are mainly protons, while for 40 TeV $<$ $E_{o}$ $<$ 100 TeV they include both $p$ and $\\alpha$ particles.\\\\ ", + "conclusions": "A measurement of the lateral distribution of Cherenkov light in EAS in the energy range 20 $\\div$ 120 TeV has been performed at the Gran Sasso Laboratories by the EAS-TOP and MACRO arrays. The EAS and its geometry are selected through the muon detected deep underground by MACRO (E$_{\\mu}$ $>$ 1.3 TeV). The measurements are performed by means of the Cherenkov light detector of EAS-TOP at Campo Imperatore (2000 m a.s.l.). The measurement is compared with the results of simulations based on the CORSIKA-QGSJET code. Simulated and real data show a good agreement, inside 20\\% systematic uncertainties.\\\\ The shape of the l.d.f. reflects the rate of energy release in the atmosphere, (i.e. the properties of the interaction, the primaries being dominated by the lightest components due to the TeV muon trigger requirement) while the absolute scale is mostly related to the event rate, i.e. the primary $p$ and $\\alpha$ spectra. The agreement of both of them (see fig.~\\ref{fig:jacee}) shows both the adequacy of the CORSIKA-QGSJET code in describing the cascades in this energy range and of the JACEE flux in the 20 $\\div$ 120 TeV region. The contribution of fluctuations and of the CNO component have been successively studied and they do not affect these conclusions." + }, + "0112/astro-ph0112369_arXiv.txt": { + "abstract": "We derive the observed luminosity function of virialized systems. Coupling this statistic with the mass function predicted by CDM cosmogonies we obtain the functional behavior of the mass--to--light ratio over a wide dynamical range in mass. We conclude that the mass--to--light ratio has a minimum close to the luminosity of an $L_*$ galaxy halo. We also show how to derive in a self-consistent way the X-ray--to--optical luminosity ratio for galaxy systems. As a consequence, we predict a fundamental break in the auto-similar scaling behavior of fundamental correlations in going from the group to the cluster scales. ", + "introduction": "The luminosity function of virialized systems (VS LF) is a relatively unexplored statistical tool. The VS LF can be used to probe a number of fundamental problems in cosmology, ranging from the efficiency of galaxy formation to the connection between X-ray and optical light in galaxy systems. Moreover the VS LF is more robust at the low--mass end of the clustering hierarchy than is the mass function obtained by traditional estimators such as projected velocity dispersions and X-ray temperatures. We constructed the VS LF (Marinoni, Hudson \\& Giuricin, ApJ submitted) after processing, with a weighting scheme, bound objects extracted from the Nearby Optical Galaxy (NOG) catalog (Giuricin et al. 2000). NOG is a statistically controlled, distance-limited ($cz_{LG}\\leq$ 6000 km/s) and magnitude-limited (B$\\leq$14) complete sample of more than $7000$ optical galaxies. The sample covers 2/3 (8.27 sr) of the sky ($|b|>20^{\\circ}$), a volume of $1.41 \\times 10^{6}\\hmpcinvcub$ and has a redshift completeness of 98\\% (see Marinoni 2001). The NOG sample spans a fair volume of the Universe, but better samples the nearby cosmological space than previous all-sky catalogues. In Fig. 1, the NOG galaxy distribution is shown using 3D isodensity surfaces of low-density contrast ($\\delta=1.5$) obtained by smoothing with a Gaussian filter of short smoothing length ($r_s = 200$ km/s). In this way we show the richness of details recovered also on small scales. In Fig. 1, we have graphically summarized the performance of the algorithms developed in order to reconstruct, in an self-consistent way, galaxy systems and pseudo real-space positions (Marinoni et al. 1998; Marinoni et al. 1999; Giuricin et al. 2000). Most of the NOG galaxies are found to be members of galaxy binaries (which comprise $\\sim$15\\% of galaxies) or groups with at least three members ($\\sim$45\\% of galaxies). About 40\\% of the galaxies are left ungrouped (isolated galaxies). ", + "conclusions": "" + }, + "0112/astro-ph0112125_arXiv.txt": { + "abstract": "We report on our search for the optical counterparts of the Southern Hemisphere anomalous X-ray pulsar 1E1048.1-5937 and the radio-quiet neutron stars in supernova remnants Puppis A, RCW 103, and PKS 1209-52. The observations were carried out with the new MIT/CfA MagIC camera on the Magellan-I 6.5 m telescope in Chile. We present deep multiband optical images of the X-ray error circles for each of these targets and discuss the resulting candidates and limits. ", + "introduction": "{\\bf The anomalous X-ray pulsars (AXPs)} are a group of X-ray pulsars whose spin periods fall in a narrow range ($\\sim 6-12$ s), whose X-ray spectra are very soft, and which show no evidence that they accrete from a binary companion (see Mereghetti 1999 for a recent review). These objects may be isolated neutron stars with extremely strong ($\\sim 10^{14}$ G) surface magnetic fields, or they may be accreting from a ``fallback'' accretion disk. Optical measurements could potentially help discriminate between these models. An optical counterpart to one AXP, 4U~0142+61, has recently been identified and shown to have peculiar optical colors (Hulleman et al. 2000).\\\\ {\\bf The radio-quiet neutron stars (RQNSs)} are a group of compact X-ray sources found near the center of young supernova remnants. Their X-ray spectra are roughly consistent with young, cooling neutron stars, but they show no evidence for the non-thermal emission associated with ``classical'' young pulsars like the Crab (see Brazier \\& Johnston 1999 for a review). The X-ray spectral properties of the RQNSs and the AXPs are similar (see, e.g., Chakrabarty et al. 2001). Below in Table 1, the general properties of the three RQNSs as our targets in the southern sky are listed. \\begin{table} \\caption{Radio-Quiet Neutron Stars} \\begin{tabular}{clcccc} \\hline & & $d$ & Age & $kT_{bb}$ & \\\\ Source & SNR & (kpc) & ($10^3$ yr) & (keV) & Refs \\\\ \\hline 1E 0820--4247 & Pup A & 2.0 & 3.7 & 0.28 & 1-3 \\\\ 1E 1614--5055 & RCW 103 & 3.3 & 1-3 & 0.56 & 4-6\\\\ 1E 1207--5209 & PKS 1209--52 & 1.5 & 7 & 0.25 & 7-9 \\\\ \\hline \\multicolumn{6}{l}{\\tiny {\\bf References.} -- (1) Petre et al. 1982. (2) Petre, Becker, \\& Winkler 1996. (3) Pavlov et al. 1999.}\\\\ \\multicolumn{6}{l}{\\tiny (4) Tuohy \\& Garmire 1980. (5) Caswell et al. 1975. (6) Gotthelf, Petre, \\& Hwang 1997. (7) Helfand \\& Becker 1984.}\\\\ \\multicolumn{6}{l}{\\tiny (8) Bignami, Caraveo, \\& Mereghetti 1992. (9) Mereghetti, Bignami, \\& Caraveo 1996.}\\\\ \\end{tabular} \\end{table} ", + "conclusions": "" + }, + "0112/astro-ph0112313_arXiv.txt": { + "abstract": "We present Giant Metrewave Radio Telescope (GMRT) observations of redshifted 21-cm absorption from the $z=0.437$ metal line absorption system towards PKS~1243-072. HI absorption is clearly detected; the absorption profile has a velocity spread of $\\sim 20$~km/s. Detection of 21-cm absorption indicates that the absorber has an HI column density large enough to be classified as a damped Lyman-$\\alpha$ system. Follow up ground based optical imaging and spectroscopy allow us to identify the absorber with an $L \\sim L^\\star$ galaxy at an impact parameter of $\\sim 9.8$~kpc from the line of sight to the QSO. The absorbing galaxy is unusual in that it has bright emission lines. On the basis of the optical spectrum we are unable to uniquely classify the galaxy since its emission line ratios lie in the transition region between starburst and Seyfert~II type spectra. ", + "introduction": "\\label{intro} Absorption lines seen in the spectra of distant quasars serve as excellent probes of intervening systems along the QSO line of sight. Of these, the highest HI column density systems, the so-called damped Lyman-$\\alpha$ absorbers (DLAs), are of particular interest as they form the major repository of neutral gas at high redshifts. By studying the evolution of these DLAs, one can observationally determine the evolution of neutral gas in the universe. The connection between the evolution of the neutral gas content and the average star formation rate in galaxies, however, remains unclear. The HI mass in DLAs at $z \\sim 3$ has been found to be comparable to the stellar mass in galaxies at $z=0$, consistent with the idea that the gas in the absorbers has been converted into stars in the intervening period (\\cite{storrie96}). As such, this makes DLAs logical candidates for the precursors of modern-day spiral galaxies (\\cite{wolfe88}). However, the deduced gas mass in DLAs also depends on the assumed cosmological parameters; in fact, both Storrie-Lombardi \\& Wolfe (2000) and Peroux et al. (2001) point out that, in the currently favoured $\\Omega_\\Lambda=0.7, \\Omega_{\\rm M}=0.3$, $H_0 = 65$~\\kms~Mpc$^{-1}$ cosmology, the estimated gas mass in DLAs at high redshift is, in fact, {\\it less} than the mass in stars at $z=0$ (albeit only at the 1~$\\sigma$ level). Interestingly, a recent Hubble Space Telescope survey for DLAs in a sample of MgII absorbers indicates that the neutral gas content in DLAs at low redshift is comparable to that at high $z$, and is, in fact, quite consistent with a scenario in which the HI has {\\it not} been converted into stars (\\cite{rt2000}). The latter is, of course, a ``biased'' survey, since the absorbers were pre-selected on the basis of their MgII absorption; the effects of this bias are unclear. Regardless of the connection between the evolution of the neutral gas density and star formation, it remains true that the study of DLAs is currently the only observational means by which one can trace the evolution of cold neutral gas in the universe. Besides the above, the typical size and structure of damped systems has also been an issue of much controversy. Proposed models for DLAs range from large, rapidly rotating proto-disks (e.g. \\cite{prochaska97}) to small, merging sub-galactic systems (e.g. \\cite{haehnelt98}). Locally, however, 21-cm emission studies indicate that spiral galaxies are the predominant contributors to the HI mass (\\cite{rao93}); one would thus expect at least the low redshift DLAs to be primarily spiral disks. However, one of the puzzling outcomes of optical imaging of low-$z$ DLAs is that such systems appear to be associated with a wide variety of galaxy types, with only a few systems originating in luminous spirals (\\cite{lebrun97,turnshek00,turnshek01,cohen01,bowen01}). Further, spectroscopic studies indicate that DLAs show very weak (if any) evolution in their metallicity with redshift and also do not show the expected $\\alpha$/Fe enrichment pattern expected for spiral galaxies (\\cite{pettini99,centurion2000}). Of course, these results could well stem from selection biases in present samples of DLAs arising from, for example, issues like dust obscuration. Such issues can be addressed by detailed imaging and spectroscopic observations of individual DLAs; such studies are, however, only possible at fairly low redshifts. Unfortunately, there are very few damped systems known at low redshifts as their identification requires UV spectra from space-based telescopes. However, all extra-galactic 21-cm absorbers for which the Lyman-$\\alpha$ line has also been observed have HI column densities $N_{\\rm HI} > 2 \\times 10^{20}$ \\cm, i.e. are classically damped. A detection of 21-cm absorption towards a radio-loud quasar can thus be used as a criterion for the identification of a damped system; this can then be followed up by optical/UV studies to identify the absorber. In this letter, we describe a search for 21-cm absorption at $z = 0.436$ towards the radio-loud quasar PKS~1243-072, using the Giant Metrewave Radio Telescope (GMRT). The quasar emission redshift is $z_{em}=1.286$ (\\cite{wilkes83}). Multiple strong low ionization absorption lines (MgII~$\\lambda{2798}$, FeII~$\\lambda{2600}, \\lambda{2587}, \\lambda{2383}$) have been detected at $z_{abs}=0.436$ towards the QSO (\\cite{wright79}). The $z=0.436$ absorber was part of the MgII-selected sample searched for 21-cm absorption by Lane (2000), with the Westerbork Synthesis Radio Telescope, and was classified as a candidate 21-cm absorber on the basis of these observations. Our fresh GMRT observations have resulted in the confirmed detection of 21-cm absorption in this system. As discussed above, this implies that the absorber fits the classical definition of a DLA. We have also carried out R and I band imaging studies of the system, as well as optical spectroscopy, resulting in the identification of the absorber with an $L \\sim L^\\star$ galaxy. ", + "conclusions": "\\subsection{HI Column Density} The 21~cm optical depth, $\\tau_{21}$, of an optically thin, homogeneous cloud is related to the column density of the absorbing gas $N_{\\rm HI}$ and the spin temperature $T_{\\rm s}$ by the expression (e.g. \\cite{rohlfs86}) \\begin{equation} \\label{eqn:tspin} N_{\\rm HI} = { 1.823\\times10^{18} T_{\\rm s} \\over f} \\int \\tau_{21} \\mathrm{d} V \\; , \\end{equation} \\noi where $f$ is the covering factor of the absorber. In the above equation, $N_{\\rm HI} $ is in cm$^{-2}$, $T_{\\rm s}$ in K and $\\mathrm{d}V$ in km s$^{-1}$. For a multi-phase absorber the spin temperature derived using the above expression is the column density weighted harmonic mean of the spin temperatures of the individual phases. In the case of PKS~1243-072, VLBA maps at 2.3 and 8.4 GHz (see the Radio Reference Frame Image Database of the United States Naval Observatory) show that the entire flux is contained within $\\sim$~20 milliarcseconds. The extremely small size of the background source makes it highly likely that the covering factor $f$ is close to unity. This implies an HI column density $N_{\\rm HI} = 3 \\times 10^{20} (T_{\\rm s}/200 {\\rm K})$~\\cm; given that all known DLAs have $T_{\\rm s} \\ge 150$~K (\\cite{chengalur00}), it is quite likely that the absorber, it is quite likely that the absorber has $N_{\\rm HI} > 2 \\times 10^{20}$~\\cm (i.e. fits the classical definition of a DLA) even if it has a low spin temperature. Of course, a spin temperature $\\ga 1000$~K, more typical of DLAs (\\cite{chengalur00}), would imply a much higher column density, $N_{\\rm HI} \\ga 10^{21}$~\\cm. However, the high luminosity of the absorber makes it likely that its spin temperature is low; this is discussed in more detail in the next section. Observations of the Lyman-$\\alpha$ line using the Hubble Space Telescope (HST) will provide a direct estimate of the column density and thus, of the spin temperature. \\subsection{The nature of the Absorber} The rather strong optical emission lines seen in the absorber are rather unusual. Further, the absorber is only slightly more extended than the quasar image (Q$_{\\rm FWHM}$ = 0.96 arcsec and A$_{\\rm FWHM}$ = 1.08 arcsec for the QSO and object~A on the R-band image, and 0.71 and 0.86~arcsec respectively, on the I-band image). In fact, object~A appears more {\\em point-like} than most of the other galaxies in the field. The emission lines could be either due to star-burst activity or due to the presence of an active galactic nucleus (AGN). It is possible to distinguish between these two scenarios on the basis of the ratios of the strengths of certain emission lines. Dessauges-Zavadsky et al. (2001; hereafter DZ01) present template spectra of different emission line galaxies and also discuss their own and earlier classification schemes (\\cite{rola97,tresse96}). A visual inspection of the template spectra clearly shows that the object~A is either a Seyfert~II system or a starburst galaxy. Next, the spectrum of object~A has the following line ratios $$ [O III] \\lambda{5007} / H\\beta \\sim 6.2 $$ $$ [O II] \\lambda{3727} / H\\beta \\sim 3 $$ $$ [O III] \\lambda{4959} / H\\beta \\sim 2.2 $$ \\noindent which can be used as diagnostics for the purpose of classification (\\cite{rola97,tresse96}). A comparison of these values with Fig. 7 of DZ01 shows that object~A falls in between the range of ratios obtained in typical HII galaxies and Seyfert~II type systems. DZ01 also presented a new diagnostic, based on a comparison between the quantity $R_{23} \\equiv ([O II] \\lambda{3727} + [O III] \\lambda{4959} + [O III] \\lambda{5007})/H_\\beta$ and the ratio $[OI] \\lambda{6300}/H\\alpha$. We unfortunately do not have a measurement of either $[OI] \\lambda{6300}$ or $H\\alpha$ in object~A; however, we do estimate $\\log{[R_{23}]}= 1.06$ for this system. DZ01 noted that 87~\\% of all Seyfert~II galaxies had $\\log{[R_{23}]} > 1.1$ and, an inspection of Fig. 8 in DZ01 shows that all Seyfert~II systems lie above $\\log{[R_{23}]} = 1.05$. Again, object~A lies close to the border separating Seyfert~II and starburst galaxies. We note that all the above diagnostics tend to place object~A marginally amongst the starburst systems. Given the lack of a measured $[OI]/H\\alpha$ ratio and the problem with spectrophotmetric calibration mentioned earlier, we are unable to conclusively distinguish between the two possibilities. We plan to carry out observations of the $[OI] \\lambda{6300}$ and $H\\alpha$ lines from this system, which should be useful in resolving the issue. The 2.2 arcsec separation between the quasar and the absorber corresponds to a linear separation of 9.8 kpc between their lines of sight at the redshift of the absorber (assuming a flat FRW Universe, with $H_0 = 75$~\\kms~Mpc$^{-1}$). Although we cannot rule out the possibility that the absorbing galaxy is not object~A, but some fainter companion galaxy, the small projected separation between object~A and the QSO makes it likely that the absorption arises in object~A itself. We note that there is a faint object 1'' north of A and 2'' east of the quasar, barely visible in Fig.~\\ref{fig:image}. This system is about 2 magnitudes fainter than object~A and considerably more diffuse. It is unclear whether this object is in the vicinity of the QSO or a companion to object~A or, indeed, an interloper not associated with either system. Our optical photometry shows that object~A has $L \\sim L^\\star$. This is consistent with the results of Rao \\& Briggs (1993) who used a survey of the HI content of $z=0$ optically bright galaxies, to conclude that the cross section for DLA absorption peaks at this luminosity. Note, however, that Rosenberg \\& Schneider (2001) argue that a substantial contribution to the DLA cross-section is provided by optically faint galaxies, based on a blind 21-cm survey at $z=0$. The latter is also consistent with optical searches for the counterparts of low redshift DLAs, which have shown that the absorbers arise in galaxies with a wide range of luminosities. Chengalur \\& Kanekar (2000) found that low spin temperatures ($T_{\\rm s} \\la 300$~K) were obtained in the few cases where the absorber was identified to be a spiral galaxy; such temperatures are typical of the Milky Way and local spirals (see also Kanekar \\& Chengalur (2001)). However, the majority of DLAs were found to have far higher spin temperatures, $T_{\\rm s} \\ga 1000$~K. Higher $T_{\\rm s}$ values are to be expected in smaller systems like dwarf galaxies, whose low metallicities and pressures are not conducive to the formation of the cold phase of HI (\\cite{wolfire95}); such systems hence have a higher fraction of warm gas as compared to normal spirals, and therefore, a high spin temperature. On the other hand, bright galaxies tend to have high masses, and hence both higher metallicities and central pressures, contributing to the formation of the cold phase of neutral hydrogen. The high luminosity of object~A thus indicates that it is likely to have a low spin temperature ($\\la 300$~K) and hence, a relatively low column density $N_{\\rm HI} \\la 5 \\times 10^{20}$~\\cm. It would be interesting to test this conjecture by means of HST observations in the Lyman-$\\alpha$ line, as well as to directly determine the metallicity of the absorber through high resolution absorption studies. In a subsequent paper, we plan to compare the metallicity as computed from such a high resolution absorption spectrum to the metallicity measured from the emission lines. \\begin{acknowledgement} The radio observations presented in this paper would not have been possible without the many years of dedicated effort put in by the GMRT staff to build the telescope. The GMRT is run by the National Centre for Radio Astrophysics of the Tata Institute of Fundamental Research. The optical observations were carried out using European Southern Observatory facilities at La Silla Observatory, Chile. This research has made use of the United States Naval Observatory (USNO) Radio Reference Frame Image Database (RRFID). This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration. \\end{acknowledgement} {\\bf Note added in proof:} 21cm absorption from 1243-072 has recently been independently detected by Lane et al. (2001) using the WSRT." + }, + "0112/astro-ph0112063_arXiv.txt": { + "abstract": "{The ISOPHOT instrument aboard ISO has been used to observe extended FIR emission of six Abell clusters. Strip scanning measurements with crossing position angles centered on the clusters were carried out at 120\\,$\\mathrm{\\mu m}$ and 180\\,$\\mathrm{\\mu m}$. The raw profiles of the $\\mathrm{I_{120 \\mu m} / I_{180 \\mu m}}$ surface brightness ratio including zodiacal light show a bump towards Abell 1656 (Coma), dips towards Abell 262 and Abell 2670, and are without clear structure towards Abell 400, Abell 496, and Abell 4038. After subtraction of the zodiacal light, the bump towards Abell 1656 is still present, while the dips towards Abell 262 and Abell 2670 are no longer noticable. This indicates a localized excess of emitting material outside the Galaxy towards Abell 1656 with properties different from the galactic foreground cirrus, while the behavior in Abell 262 and Abell 2670 can be reconciled with galactic cirrus structures localized on the line-of-sight to these clusters. The excess of $\\approx$ 0.2 MJy/sr seen at 120\\,$\\mathrm{\\mu m}$ towards Abell 1656 (Coma) is interpreted as being due to thermal emission from intracluster dust distributed in the hot X-ray emitting intracluster medium. The integrated excess flux within the central region of 10\\,$\\arcmin$ -- 15\\,$\\arcmin$ diameter is $\\approx$ 2.8 Jy. Since the dust temperature is poorly constrained, only a rough estimate of the associated dust mass of $\\mathrm{M_{D} \\approx 10^{7}\\,M_{\\sun}}$ can be derived. The associated visual extinction is negligible ($\\mathrm{A_{V} \\ll 0.1 mag}$) and much smaller than claimed from optical observations. No evidence is found for intracluster dust in the other five clusters observed. The absence of any signature for intracluster dust in five clusters and the rather low inferred dust mass in Abell 1656 indicates that intracluster dust is likely not responsible for the excess X-ray absorption seen in cooling flow clusters. These observations thereby represent a further unsuccessful attempt in detecting the presumed final stage of the cooling flow material, in accord with quite a number of previous studies in other wavelengths regions. Finally, the observed dimming of the high-redshift supernovae is unlikely be attributable to extinction caused by dust in the intracluster or even a presumed intercluster medium. ", + "introduction": " ", + "conclusions": "" + }, + "0112/astro-ph0112549_arXiv.txt": { + "abstract": "We propose that some of the high-latitude unidentified EGRET $\\gamma$-ray sources could be the result of gravitational lensing amplification of the innermost regions of distant, faint, active galactic nuclei. These objects have $\\gamma$-ray emitting regions small enough as to be affected by microlensing of stars in interposed galaxies. We compute the gravitational amplification taking into account effects of the host galaxy of the lens and prove that, whereas the innermost $\\gamma$-ray regions can be magnified up to thousand times, there is no amplification at radio frequencies, leading to the observed absence of strong counterparts. Some new effects in the spectral evolution of a gravitational microlensed $\\gamma$-ray AGN are predicted. Within a reasonable range of lensing parameters, and/or types of sources, both variable and non-variable EGRET detections at high latitudes can be explained by microlensing. The same phenomenon could also have an important incidence among the future GLAST detections at high-latitudes. ", + "introduction": "The Third EGRET Catalog of $\\gamma$-ray sources includes observations carried out between April 22, 1991 and October 3, 1995, and lists 271 point-like detections (Hartman et al. 1999). About two thirds of them have no conclusive counterparts at lower frequencies, and are dubbed unidentified. These unidentified $\\gamma$-ray sources can be divided in two broad groups. The first one, at low latitudes, is probably related to several galactic populations such as radio-quiet pulsars, interacting supernova remnants, early-type stars with strong stellar winds, X-ray binaries, etc (see Romero 2001 and references therein). The second group of unidentified sources is formed by mid- and high-latitude detections. Gehrels et al. (2000) have shown --in a model independent way-- that the mid-latitude sources are different from the bright population of unidentified sources along the galactic plane. Some of the mid-latitude detections ($5^{\\circ}<|b|<30^{\\circ}$) are thought to be associated with the Gould Belt (Grenier 2000, Gehrels et al. 2000), a starburst region lying at $\\sim600$ pc from Earth. A few other sources, at higher latitudes, might be the result of electrons being accelerated at the shock waves of forming clusters of galaxies (Totani et al. 2001). However, for most of the high-latitude unidentified sources (HL-UnidS), no other explanation seems to be available than they are AGNs yet unnoticed at lower energies (Reimer \\& Thompson 2001). A $\\gamma$-ray AGN population consisting of 66 members has already been detected by EGRET (Hartman et al. 1999). Photon spectra and variability indices of the HL-UnidS match well those of known $\\gamma$-ray AGNs. In Figure 1, we show the corresponding distributions, with the AGN-histogram adapted from Torres et al. (2001a). The variability criterion used here is the $I$-scheme, but other variability criteria yield similar results (Torres et al. 2001b). Population studies have already remarked that part of the sample of HL-UnidS is consistent with an isotropic population, a fact that also supports an extragalactic origin (\\\"Ozel \\& Thompson 1996). All identified $\\gamma$-ray AGNs are also strong radio sources with flat spectrum, as expected from synchrotron jet-like sources where the $\\gamma$-ray flux is the result of inverse Compton scattering (Mattox et al. 1997). We could ask, then, why the HL-UnidS are not detected at lower frequencies, particularly in the radio band, if they are also AGNs? Here we propose that some of these sources are the result of gravitational lensing amplification of background, high-redshift, active galactic nuclei; blazars whose $\\gamma$-ray emitting regions are small enough as to be affected by microlensing by stars in interposed galaxies. ", + "conclusions": "In summary, we have found that gravitational microlensing of the innermost regions of distant AGNs can produce unidentified sources compatible with those observed at high galactic latitudes. While large amplification factors are obtained for $\\gamma$-rays, a negligible magnification results in the radio band. In the case of the optical emission, if it is the same particle population giving rise to gamma-rays and optical emission then we can also expect magnification of the optical luminosity. The gamma-ray spectral evolution provides a specific signature for the microlensing events that can be used to differentiate this from other kind of phenomena. Higher (lensing) optical depth, or the presence of shear, will lead to a diversity of light curves. We remark that we have used the Chang and Refsdal (1979) model as the gravitational lensing scenario, but that galaxies with denser cores would require a more detailed treatment. This would necessarily include the study of caustic patterns of all stars at a time in order to get the light curves. We shall explore the lensing model in much more detail in a subsequent publication (Eiroa et al. 2002, in preparation).\\\\" + }, + "0112/astro-ph0112255_arXiv.txt": { + "abstract": "Cold Dark Matter (CDM) has become the standard modern theory of cosmological structure formation. Its predictions appear to be in good agreement with data on large scales, and it naturally accounts for many properties of galaxies. But despite its many successes, there has been concern about CDM on small scales because of the possible contradiction between the linearly rising rotation curves observed in some dark-matter-dominated galaxies vs.~the $1/r$ density cusps at the centers of simulated CDM halos. Other CDM issues on small scales include the very large number of small satellite halos in simulations, far more than the number of small galaxies observed locally, and problems concerning the angular momentum of the baryons in dark matter halos. The latest data and simulations have lessened, although not entirely resolved, these concerns. Meanwhile, the main alternatives to CDM that have been considered to solve these problems, self-interacting dark matter (SIDM) and warm dark matter (WDM), have been found to have serious drawbacks.\\footnote{This paper is a significantly updated revision of~\\protect\\cite{cosmo2000}.} ", + "introduction": "Sometimes a theory is proposed in relatively early stages of the development of a scientific field, and this theory turns out to be not only a useful paradigm for the further development of the field --- it also survives confrontation with a vast amount of data, and becomes accepted as the standard theory. This happened with General Relativity \\cite{will}, and it seems to be happening now with general relativistic cosmology. It appears that the universe on the largest scales can indeed be described by three numbers: \\begin{itemize} \\item $H_0 \\equiv 100 h \\kmsMpc$, the Hubble parameter (expansion rate of the universe) at the present epoch, \\item $\\Omega_m \\equiv \\rho/\\rho_c$, the density of matter $\\rho$ in units of critical density $\\rho_c \\equiv 3 H_0^2 (8\\pi G)^{-1} = 2.78 \\times 10^{11} h^2 M_\\odot$ Mpc$^{-3}$, and \\item $\\Omega_\\Lambda \\equiv \\Lambda (3H_0^2)^{-1}$, the corresponding quantity for the cosmological constant. \\end{itemize} The currently measured values of these and other key parameters are summarized in the Table below. It remains to be seen whether the ``dark energy'' represented by the cosmological constant $\\Lambda$ is really constant, or is perhaps instead a consequence of the dynamics of some fundamental field as in ``quintessence'' theories \\cite{quintessence}. In particle physics, the first unified theory of the weak and electromagnetic interactions \\cite{so3} had as its fundamental bosons just the carriers of the charged weak interactions $W^+$, $W^-$, and the photon $\\gamma$. The next such theory \\cite{su2xu1} had a slightly more complicated pattern of gauge bosons --- a triplet plus a singlet, out of which came not only $W^+$, $W^-$, and $\\gamma$, but also the neutral weak boson $Z^0$, and correspondingly an extra free parameter, the ``Weinberg angle.'' It was of course this latter SU(2)$\\times$U(1) theory which has now become part of the Standard Model of particle physics. During the early 1970s, however, when the experimental data were just becoming available and some of the data appeared to contradict the SU(2)$\\times$U(1) theory, many other more complicated theories were proposed, even by Weinberg \\cite{su3xsu3}, but all these more complicated theories ultimately fell by the wayside. The development of theories of dark matter may follow a similar pattern. By the late 1970s it was becoming clear both that a great deal of dark matter exists \\cite{fabergallagher} and that the cosmic microwave background (CMB) fluctuation amplitude is smaller than that predicted in a baryonic universe. The first nonbaryonic dark matter candidate to be investigated in detail was light neutrinos --- what we now call ``hot dark matter'' (HDM). This dark matter is called ``hot'' because at one year after the big bang, when the horizon first encompassed the amount of matter in a large galaxy like our own (about $10^{12} M_\\odot$) and the temperature was about 1 keV \\cite{varenna}, neutrinos with masses in the eV range would have been highly relativistic. It is hardly surprising that HDM was worked out first. Neutrinos were known to exist, after all, and an experiment in Moscow that had measured a mass for the electron neutrino $m(\\nu_e) \\approx 20$ eV (corresponding to $\\Omega_m \\approx 1$ if $h$ were as small as $\\sim0.5$, since $\\Omega_\\nu = m(\\nu_e) (92 h^2 {\\rm eV})^{-1}$) had motivated especially Zel'dovich and his colleagues to work out the implications of HDM with a Zel'dovich spectrum ($P_p(k)=Ak^n$ with $n=1$) of adiabatic primordial fluctuations. But improved experiments subsequently have only produced upper limits for $m(\\nu_e)$, currently about 3 eV \\cite{pdb}, and the predictions of the adiabatic HDM model are clearly inconsistent with the observed universe \\cite{primackgross,beamline}. Cold Dark Matter (CDM) was worked out as the problems with HDM were beginning to become clear. CDM assumes that the dark matter is mostly cold --- i.e., with negligible thermal velocities in the early universe, either because the dark matter particles are weakly interacting massive particles (WIMPs) with mass $\\sim 10^2$ GeV, or alternatively because they are produced without a thermal distribution of velocities, as is the case with axions. The CDM theory also assumes, like HDM, that the fluctuations in the dark matter have a nearly Zel'dovich spectrum of adiabatic fluctuations. Considering that the CDM model of structure formation in the universe was proposed almost twenty years ago \\cite{peeb82,pb83,bfpr}, its successes are nothing short of amazing. As I will discuss, the \\lcdm\\ variant of CDM with $\\Omega_m = 1-\\Omega_\\Lambda \\approx 0.3$ appears to be in good agreement with the available data on large scales. Issues that have arisen on smaller scales, such as the centers of dark matter halos and the numbers of small satellites, have prompted people to propose a wide variety of alternatives to CDM, such as self-interacting dark matter (SIDM) \\cite{sidm}. It remains to be seen whether such alternative theories with extra parameters actually turn out to be in better agreement with data. As I will discuss below, it now appears that SIDM is probably ruled out, while the small-scale predictions of CDM may be in better agreement with the latest data than appeared to be the case as recently as a year ago. In the next section I will briefly review the current observations and the successes of \\lcdm\\ on large scales, and then I will discuss the possible problems on small scales. ", + "conclusions": "" + }, + "0112/astro-ph0112038_arXiv.txt": { + "abstract": "{In this work we investigate the gravitationally lensed system B1422+231. High-quality VLBI image positions, fluxes and shapes as well as an optical HST lens galaxy position are used. First, two simple and smooth models for the lens galaxy are applied to fit observed image positions and fluxes; no even remotely acceptable model was found. Such models also do not accurately reproduce the image shapes. In order to fit the data successfully, mass substructure has to be added to the lens, and its level is estimated. To explore expectations about the level of substructure in galaxies and its influence on strong lensing, N-body simulation results of a model galaxy are employed. By using the mass distribution of this model galaxy as a lens, synthetic data sets of different four image system configurations are generated and simple lens models are again applied to fit them. The difficulties in fitting these lens systems turn out to be similar to the case of some real gravitationally lensed systems, thus possibly providing evidence for the presence and strong influence of substructure in the primary lens galaxy. ", + "introduction": "} Gravitational lens systems with multiply imaged quasars are an excellent tool for studying the properties of distant galaxies. In particular, they provide the most accurate mass measures for the lensing galaxy. Besides the mass profiles, one can also gain information about evolution \\citep{ko00} and extinction laws \\citep{fa99}. Strong lensing is also a very promising and robust tool to measure the Hubble constant \\citep{re64}. The success of this method, however, depends strongly on how well the mass model is constrained. It turns out that image positions can be fit quite accurately with simple, smooth elliptical models \\citep{ke97}. Since the number of observational constraints from image positions is small, one wants to include the flux information. The optical fluxes, however, should not be used, as they might be affected by microlensing and/or dust obscuration \\citep{ch79}. Radio fluxes, on the other hand, can provide further constraints in lens modelling. Fitting the fluxes very accurately turns out to be difficult in many lens systems. Models for MG0414+0534 \\citep{fa97,ke97}, PG1115+080 \\citep{ke97b} and B1422+231 (e.g. \\citealt{ko94b}) all show the same failure, namely the observed flux ratios are very different from what one would expect for the image configurations from a smooth model. In particular, the gravitational lens system B1422+231 was explored in detail, and as was first mentioned by \\citet{ma98}, mass substructure in the lens galaxy might provide an explanation for the failures in flux modelling. A question arises as to whether there is something special about these quadruple systems, or does the discrepancy simply arise due to the fact that our smooth models are oversimplified? In other words, we are asking how ``well'' (for the purpose of strong lensing) can the smooth models used to fit the data represent a real galaxy? N-body simulation data can provide a realistic description of a galaxy mass distribution, thus giving a possibility to probe its effect on strong lensing. In this paper, we will study the influence of the substructure on the lens system B1422+231, using VLBI radio measurements of the system by \\citet{pa99}. In Sec.~\\ref{sc:1422} we give a description of the lens system and data used. The method is outlined in Sect.~\\ref{sc:lm} and the results on fitting the system with smooth mass model are presented. In Sect.~\\ref{sc:mss} the model accounting for the substructure is presented and in Sect.~\\ref{sc:ell} the deconvolved image shape information is added to the fit. Sect.~\\ref{sc:nbody} gives the description of the method used to investigate lensing by an N-body simulated galaxy. We describe how we obtained synthetic data of four image systems, and the results of fitting such systems with lens models are presented. Finally we draw some conclusions in Sect.~\\ref{sc:concl}. This work is an abbreviated version of \\citet{diploma}. In the course of writing this paper, several related papers on substructure of lens galaxies have been submitted \\citep{me01,ch01,da01,me01b,ke01}. In \\citet{me01} and \\citet{ch01}, the authors use semi-analytical descriptions to account for mass clumps typical for N-body simulations in a statistical fashion. In addition, \\citet{ch01} tests his prediction on two known systems with four images. \\citet{da01}, \\citet{me01b}, and \\citet{ke01} predict the properties of substructure needed to constrain lens systems and compare their results with the CDM predictions. In addition, \\citet{ke01} accounts for the difference between radio and optical fluxes by investigating radio and optical sources of different sizes. The present work is differs from the afore mentioned in that slightly different models for the lens in B1422+231 system are used and that we also include image shape constraints in the fit. Further, we investigate four image systems generated by using a CDM N-body simulated galaxy, rather than an analytic approximation for the statistics of mass-substructure. ", + "conclusions": "} In this work we have investigated the influence of substructure in the gravitationally lensed system B1422+231. While it is intuitively clear that a lens galaxy is not a smooth entity, we have tried to investigate how deviation from a smooth model can influence lensing phenomena, especially the image flux ratios. We have used two different smooth models for the lensing galaxy (SIE+SH and NIE+SH), and both failed very badly in fitting the image fluxes (we got $\\chi^2=130$ with 2 degrees of freedom). The use of models with substructure requires additional observational constraints. Therefore, we used deconvolved image shapes as constraints. We get a significant improvement of the fit compared to the smooth model. However, the way the substructure is introduced is oversimplified, thus we should not be surprised that the resulting $\\chi^2$ is still high. For the model with a single perturber we got $\\chi^2=19$ for 6 degrees of freedom, and with two perturbers we had $\\chi^2=15$ for 4, while the model without substructure (where deconvolved image shapes were included) gives $\\chi^2=140$ for 8 degrees of freedom. Up to now we have not considered the possibility that microlensing plays a role for the radio fluxes. \\citet{kop00} claim that they have detected microlensing in the multiply-imaged radio source B1600+434. Microlensing is a very tempting explanation for difficulties in fitting the fluxes for it can also explain why the $8.4 \\: {\\rm GHz}$ A:B flux ratio has changed from 0.97 in 1991 \\citep{pa92} to 0.93 in 1997 \\citep{pa99}. This has a consequence that again speaks in favour of substructure, since the presence of radio microlensing indicates that there is a significant number of compact objects in the lens galaxy halo. N-body simulation data of a model galaxy provides a test for the influence of mass-substructure in strong gravitational lensing. When we generated data of four image systems with the simulated galaxy we again experienced difficulties. We have tried to fit image positions and fluxes and failed to obtaining a model that fits well. From these experiments we can conclude that the level of substructure obtained from this particular N-body simulated galaxy can cause the same difficulties as experienced in some of the real gravitationally lensed systems. In order to obtain stronger conclusions one would have to investigate more realisations of simulated galaxies, also at different redshifts. However, the N-body simulation results indicate that substructure plays an important role in strong lensing. Also, modelling B1422+231 shows that the fluxes of more than one image are probably affected by the lens substructure. With the conclusions one can draw from lens modelling we can not give a direct answer about the properties of mass substructure in B1422+231. One should therefore avoid using the flux constraints directly; they should, rather, be treated in statistical manner, e.g. in a way suggested by \\citet{ma98}. Fortunately, the perturbations on the scales we are dealing with here do not influence the image positions significantly, and play even less of a role for the time delay \\citep{ma98}. Strong lensing thus remains one of the best tools to constrain the Hubble constant." + }, + "0112/astro-ph0112381_arXiv.txt": { + "abstract": "This article briefly summarizes the increasingly precise observational estimates of the cosmological parameters. After three years on the stump, the $\\Lambda$CDM model is still the leading candidate. Although the Universe is expanding, our picture of it is coming together. ", + "introduction": "If your model of the Universe is a turtle, you want to know how big the turtle is, how old the turtle is, where the turtle came from and, in some obscure animistic models, how fast the turtle is expanding. Cosmological parameters are the observable quantities that most cosmologists think are important. In the context of general relativity and the hot big bang model, cosmological parameters are the numbers that, when inserted into the Friedmann equation, \\begin{equation} H^{2} = H_{o}^{2}\\left[\\ol + \\ok \\;a^{-2} + \\om \\;a^{-3} + \\orel \\;a^{-4}\\right], \\label{eq:Friedmann} \\end{equation} best describe our particular observable Universe. The expansion is parametrized by Hubble's constant, $H_{o} = \\dot{a}/a$, where $a$ is the scale factor of the Universe. Observational estimates of the parameters in this equation, $H_{o}$, $\\ol$, $\\ok$, $\\om$ and $\\orel$ (and their subcomponents) have been derived from hundreds of observations and analyses (e.g. Fig.~\\ref{f:omol}). Table 1 is my attempt to summarize this immense body of work. {\\footnotesize % \\begin{table}[!hb] \\begin{center} \\caption{Cosmological Parameters (Background)} % \\begin{tabular}{|l|l|l|l|} \\hline \\multicolumn{1}{|c}{Parameter} & \\multicolumn{1}{|c}{Estimate} & \\multicolumn{1}{|c}{Sub-Components} & \\multicolumn{1}{|c|}{References}\\\\ \\hline % cosmological constant$^{a}$ &$\\ol = 0.7 \\pm 0.1 $& &\\cite{L98} -- \\cite{deBernardis}\\\\ matter$^{b}$ &$\\om = 0.3 \\pm 0.1 $& & ` ' + \\cite{om}\\\\ \\hfill cold dark matter & &$\\oc= 0.26 \\pm 0.1 $& ` ' + \\cite{bbn}\\\\ \\hfill baryonic matter$^{c}$ & &$\\ob = 0.04 \\pm 0.01 $&\\cite{bbn}\\\\ relativistic component$^{d}$ &$ 0.01 \\lsim \\orel \\lsim 0.05$& &\\cite{Wang} \\cite{Durrer} \\cite{nu}\\\\ \\hfill neutrinos$^{e}$ & &$ 0.01 \\lsim \\on \\lsim 0.05$& ` '\\\\ \\hfill photons$^{f}$ & &$\\og=4.8^{+1.3}_{-0.9} \\times 10^{-5}$&\\cite{tcmb}\\\\ Hubble's constant$^{g}$ &$h= 0.72 \\pm 0.08 $& &\\cite{h}\\\\ age of Universe$^{h}$ &$\\t = 13.4 \\pm 1.6$ Gyr & &\\cite{L99}\\cite{Jaffe}\\cite{TZH}\\cite{Ferreras}\\\\ geometry$^{i}$ &$\\ok = 0.00 \\pm 0.06 $& &` ' + \\cite{Wang}\\cite{deBernardisNat}\\cite{Dodelson}\\\\ equation of state$^{j}$ &$w = -1.0^{+0.4} $& &\\cite{w} \\\\ deceleration parameter$^{k}$ &$q_{o} = -0.05 \\pm 0.15 $& &\\cite{sn}\\\\ CMB temperature &$T_{\\rm CMB} = 2.725 \\pm 0.001$ K& &\\cite{tcmb}\\\\ \\hline \\end{tabular}\\\\ \\end{center} \\scriptsize % \\noindent $^{a}$ $\\ol = \\frac{\\rho_{\\Lambda}}{\\rho_{crit}}$, $\\rho_{crit} = \\frac{3\\: H_{o}^{2}}{8 \\pi G}$, $\\rho_{\\Lambda} = \\frac{\\Lambda}{3\\;H_{o}^{2}}$, $^{b}$ $\\om = \\oc + \\ob$, $^{c}$ $\\obh2=0.020\\pm 0.002$ \\cite{bbn}, $^{d}$ $\\orel = \\on +\\og$, $^{e}$ $ 0.04\\; eV < m_{\\nu,\\tau} < 4.4\\; eV$ \\cite{Wang},\\cite{Durrer}, $^{f}$ $\\og =2.47 \\times 10^{-5}\\; h^{-2}\\; T_{2.725}^{4}$ \\cite{Scott}, $^{g}$ $ h = H_{o}/100 \\; km^{-1} s^{-1}Mpc^{-1}$, $^{h}$ $\\t = h^{-1}f(\\om, \\ol)$, see Fig.\\ref{f:age}, $^{i}$ $\\ok = 1 - \\otot$, $\\otot = \\ol + \\om + \\orel$, $\\ok = 0$ (flat), $>0$ (open), $<0$ (closed), thus $\\otot = 1.00 \\pm 0.06$, $^{j}$ $p = w\\rho$, $^{k}$ $q_{o} = \\om - \\ol/2$ \\end{table} } % \\clearpage \\begin{figure}[!ht] \\centering \\resizebox{8cm}{8cm}{\\includegraphics[60,80][538,608]{plotchi1.ps}} \\resizebox{8cm}{8cm}{\\includegraphics{energydensitiesxv.ps}} \\caption{Various combinations of cosmic microwave background (CMB), supernovae and other observational constraints favor the region $(\\om, \\ol) \\approx (0.3, 0.7)$ \\cite{L98} -- \\cite{Roos}. The most recent analyses, \\cite{Stompor} \\cite{deBernardis}, continue to favor this region. The composition of the Universe is thus, energy of the vacuum: $70\\% \\pm 10\\%$, matter: $30\\% \\pm 10\\%$ (cold dark matter: $26\\% \\pm 10\\%$, normal baryonic matter: $4\\% \\pm 1\\%$), with negligible energy density from photons. The neutrino energy density is poorly constrained and may be as large as the baryonic energy density. About 13\\% of the matter in the Universe is baryonic ($\\frac{\\ob}{\\om} = \\frac{0.04}{0.3} = 0.13$). % The baryons can be further divided into $3\\%$ warm invisible gas, $0.5\\%$ optically visible stars and $0.5\\%$ hot gas visible in the x-rays \\cite{Fukugita}. The $\\om$ in the top plot is equal to the sum of the $\\oc$ and $\\ob$ in the lower plot.} \\label{f:omol} \\end{figure} \\clearpage Various methods to extract cosmological parameters from cosmic microwave background (CMB) and non-CMB observations are forming an ever-tightening network of interlocking constraints. CMB observations tightly constrain $\\ok$, while type Ia supernovae observations tightly constrain $q_{o}$. Since lines of contant $\\ok$ and constant $q_{o}$ are nearly orthogonal in the $\\om-\\ol$ plane, combining these measurements optimally constrains our Universe to a small region (Fig. 1). Four years ago when $\\ol$ was assumed to be zero, the critical density, $\\ocrit = \\frac{3 H_{o}^{2}}{8\\pi G}$ {\\em was} critical -- it determined the fate of the universe -- whether it would expand forever or recollapse. Currently, the notion of critical density has lost much of its importance. That role has been usurped by $\\ol$; if $\\ol > 0$ the universe will expand forever. The upper limit on the energy density of neutrinos comes from the shape of the small scale power spectrum. If neutrinos make a significant contribution to the density, they suppress the growth of small scale structure by free-streaming out of over-densities. The CMB power spectrum is not sensitive to such suppression and is not a good way to constrain $\\on$. Hubble scholars used to be irreconcilably divided into camps described by a bimodal distribution peaking at $H_{o} = 50$ and $H_{o} = 90$. These peaks seem to have merged into a more agreeable Gaussian distribution peaking between 65 and 80 with error bars from hostile groups now overlapping. The parameters in Table 1 are not independent of each other. The elongated contours in the top plot of Fig. 1 is one example of correlation. Another example is the age of the Universe, $\\t = h^{-1}f(\\om, \\ol)$. Estimates of $h, \\ol, \\om$ can be inserted into Eq.~\\ref{eq:Friedmann}. Integration then yields the age of the Universe ($\\orel$ is negligible and $\\ok = 1- \\om - \\ol \\approx 0$). If the Universe is to make sense, independent determinations of $\\ol$, $\\om$ and $h$ and the minimum age of the Universe must be consistent with each other. This is now the case (Fig. 2). Presumably we live in a Universe which corresponds to a single point in multidimensional parameter space. Estimates of $h$ from HST Cepheids and the CMB must overlap. Deuterium and CMB determinations of $\\obh2$ should be consistent. Regions of the $\\om - \\ol$ plane favored by supernovae and CMB must overlap with each other and with other independent constraints. This is the case \\cite{L99}. The proportionality constant $w$ in the equation of state, $p = w\\rho$, is important in deciding whether the generalization of $\\ol$ into a time varying $\\ol$ (i.e. quintessence) is necessary. So far the observations seem to be favoring the simplest case $w= -1$, (i.e. pure cosmological constant) and do not call for this generalization. Just as $h$ describes the first derivative of the scale factor, the deceleration parameter $q_{o}$ describes the second derivative. The redshifts and apparent magnitudes of type Ia supernovae have been used to find that $q_{o} \\lsim 0$ -- the expansion of the Universe is accelerating. The geometry of the Universe does not seem to be like the surface of a ball ($\\ok < 0$) nor like a saddle ($\\ok > 0$) but seems to be flat ($\\ok \\approx 0$) to the precision of our current observations. \\clearpage \\begin{figure}[!h] \\centering \\includegraphics[height=15cm,width=14cm]{t2.ps} \\caption{Estimates of the age of the Universe based on Eq. 1 are plotted on the far right (`Universe') and can be compared to lower limits on the age of the Universe from age estimates of the halo and disk of our Milky Way Galaxy. Many different techniques, data sets and analysis methods were used to obtain these estimates. Figure modified from \\cite{L99}.} \\label{f:age} \\end{figure} ", + "conclusions": "" + }, + "0112/astro-ph0112348_arXiv.txt": { + "abstract": "By extrapolating to O/H = N/H = 0 the empirical correlations $Y$--O/H and $Y$--N/H defined by a relatively large sample of $\\sim$ 45 Blue Compact Dwarfs (BCDs), we have obtained a primordial $^4$Helium mass fraction $Y_{\\rm p}$= 0.2443$\\pm$0.0015 with d$Y$/d$Z$ = 2.4$\\pm$1.0. This result is in excellent agreement with the average $Y_{\\rm p}$= 0.2452$\\pm$0.0015 determined in the two most metal-deficient BCDs known, I Zw 18 ($Z_\\odot$/50) and SBS 0335--052 ($Z_\\odot$/41), where the correction for He production is smallest. The quoted error (1$\\sigma$) of $\\la$ 1\\% is statistical and does not include systematic effects. We examine various systematic effects including collisional excitation of Hydrogen lines, ionization structure and temperature fluctuation effects, and underlying stellar He {\\sc i} absorption, and conclude that combining all systematic effects, our $Y_{\\rm p}$ may be underestimated by $\\sim$ 2--4\\%. Taken at face value, our $Y_{\\rm p}$ implies a baryon-to-photon number ratio $\\eta$ = (4.7$^{+1.0}_{-0.8}$)$\\times$10$^{-10}$ and a baryon mass fraction $\\Omega_b$$h^2_{100}$ = 0.017$\\pm$0.005 (2$\\sigma$), consistent with the values obtained from deuterium and Cosmic Microwave Background measurements. Correcting $Y_{\\rm p}$ upward by 2--4\\% would make the agreement even better. ", + "introduction": "The standard hot big bang model of nucleosynthesis (SBBN) is one of the key quantitative tests of big bang cosmology, along with the Hubble expansion and the cosmic microwave background radiation. In the SBBN, four light isotopes, D, $^3$He, $^4$He and $^7$Li, were produced by nuclear reactions a few minutes after the birth of the Universe. Given the number of relativistic neutrino species and the neutron lifetime, the abundances of these light elements depend on one cosmological parameter only, the baryon-to-photon ratio $\\eta$, which in turn is directly related to the density of ordinary baryonic matter $\\Omega_b$. The ratio of any two primordial abundances, for example that of D to H gives $\\eta$, and accurate measurements of the other three light elements, for example $^4$He/H, tests SBBN. Of all light elements, the abundance of deuterium (D) is the most sensitive to the baryonic density. The primordial D abundance can be measured directly in low-metallicity absorption line systems in the spectra of high-redshift quasars. The quasar is used as a background light source, and the nearly primordial gas doing the absorbing is in the outer regions of intervening galaxies or in the intergalactic medium (the so-called Lyman $\\alpha$ clouds). Tytler and his group (see Tytler et al. 2000 for a review) have vigorously pursued this type of measurements. They have now obtained D/H measurements in the line of sight towards 4 quasars. Combining all measurements, they found all their data are consistent with a single primordial value of the D/H ratio: (D/H)$_{\\rm p}$ = 3.0$\\pm$0.4$\\times$10$^{ -5}$ (O'Meara et al. 2001). This latest value is about 10\\% lower than their previous value (D/H)$_{\\rm p}$ = 3.39$\\pm$0.25$\\times$10$^{-5}$ (Burles \\& Tytler 1998). The primordial abundance of $^3$He is also quite sensitive to the baryon density, though less than the D abundance. It has not been yet measured, mainly because low-mass stars make a lot of $^3$He, increasing its value in the interstellar medium of the Milky Way well above the primordial value. Furthermore, the amount of $^3$He destroyed in stars is unknown. Bania, Rood \\& Balser (2000) have measured an average $^3$He/H = 1.5$\\pm$0.6$\\times$10$^{-5}$ in Galactic H {\\sc ii} regions. This value represents the average in the interstellar medium of the Milky Way, but there exists no good way to extrapolate the $^3$He abundance to the primordial value. Old halo stars that formed from nearly pristine gas with very low iron abundances during the gravitational collapse of the Milky Way show approximately constant $^7$Li/H (the so called ``Spite plateau'', Spite \\& Spite 1982), implying that their $^7$Li is nearly primordial. Creation or depletion of $^7$Li may make the $^7$Li abundances of halo stars deviate from the primordial value. Creation of $^7$Li in the interstellar medium by cosmic ray spallation prior to the formation of the Milky Way has to be less than 10--20\\%, so as not to produce more Be than is observed (Ryan, Norris \\& Beers 1999). There is still considerable debate concerning the possible depletion of $^7$Li at the surface of stars. Depletion mechanisms that have been proposed include mixing due to rotation or gravity waves, mass loss in stellar winds and gravitational settling. Depending on the exact depletion mechanism, the primordial lithium abundance varies from ($^7$Li/H)$_{\\rm p}$ = (1.73$\\pm$0.21)$\\times$10$^{-10}$ (Bonifacio \\& Molaro 1997) to (2.24$\\pm$0.57)$\\times$10$^{-10}$ (Vauclair \\& Charbonnel 1998), to (3.9$\\pm$0.85)$\\times$10$^{-10}$ (Pinsonneault et al. 1999). Because of the relatively large uncertainties in the determination of the primordial abundances of $^3$He and $^7$Li, the primordial abundance of $^4$He plays a key role for deriving $\\Omega_b$ independently of D measurements, and is crucial for checking the consistency of SBBN. We discuss next how the primordial $^4$He mass fraction $Y_{\\rm p}$ is determined, and the uncertainties which enter in such a determination. ", + "conclusions": "" + }, + "0112/astro-ph0112454_arXiv.txt": { + "abstract": "We present new mathematical alternatives for explaining rotation curves of spiral galaxies in the MOND context. For given total masses, it is shown that various mathematical alternatives to MOND, while predicting flat rotation curves for large radii ($r/r_d\\gg 4$, where $r_d$ is the characteristic radius of the galactic disc), predict curves with different peculiar features for smaller radii ($0.130^{\\circ}$), many of which may be blazars. ", + "introduction": "Since the earliest days of gamma-ray astronomy, one of the foremost questions in the field has been the identity of discrete sources. In 1972-1973, SAS-2 was the first mission to detect radiation from the Vela and Crab pulsars (Fichtel et al. 1975). Launched in 1975, COS-B built on the success of SAS-2 by detecting 25 point sources. However, only four of those were identified (Swanenburg et al. 1981). As instruments became more sophisticated, the number of identified and unidentified sources has increased. The present list (Hartman et al. 1999, hereafter H99) is comprised of the 271 sources detected by the Energetic Gamma-ray Experiment Telescope (EGRET) on the late Compton Gamma-ray Observatory (CGRO) that display significant flux above 100 MeV. Of these, 169 remain unidentified. The 102 identified sources include a probable association with the radio galaxy Cen A, a solar flare, and the LMC. The remaining sources are pulsars (5), AGNs with low-confidence (27), and AGNs with high-confidence (66). The AGNs are blazars, typically flat-spectrum radio quasars (FSRQs) or BL Lac objects. There is statistical evidence that supernova remnants (Sturner \\& Dermer 1995; Esposito et al. 1996), OB associations (Kaaret \\& Cottam 1996; Romero, Benaglia, \\& Torres 1999), and objects born in the Gould Belt (Gehrels et al. 2000) are also gamma-ray emitters, but no single source type has been conclusively associated with any of these classes. Due to the relatively poor angular resolution of even the best gamma-ray telescopes, efforts to identify sources depend on factors other than spatial coincidence. Pulsars are typically found in the Galactic plane while most blazars are found at high latitude where the diffuse Galactic emission does not overwhelm the source photons. Additionally, variability over hours and days offers evidence against identification as a pulsar (Ramanamurthy et al. 1995) while blazars are known to have strongly variable gamma-ray flux (e.g. Mattox et al. 1997a). But latitude and variability studies cannot be used alone to identify individual gamma-ray sources with high confidence as pulsars or blazars; for pulsars, a clear pulse profile must be detected, while for blazars, multiwavelength studies are necessary as these sources have a distinct broad-band signature. They are variable at many frequencies, feature flat radio spectra, show variable polarization at radio and optical frequencies, display power-law spectra at X-ray and gamma-ray energies, and have moderate to large redshifts. Recently there have been a number of papers published that describe efforts to identify individual gamma-ray sources using multiwavelength analyses (e.g., Mukherjee et al. 2000; Halpern et al. 2001; Mirabel \\& Halpern 2001; Reimer et al. 2001a). In the present paper a multiwavelength study of 3EG J2006--2321 is presented, and the data indicate that this source, listed as unidentified in H99, is a blazar. In \\S\\S 2-5 the relevant data are presented; this is followed in \\S 6 by a short summary of the multifrequency data. The weak flux density of the radio counterpart and its relevance to other high-latitude unidentified EGRET sources is discussed in \\S 7. In \\S 8 conclusions and suggestions for future work are summarized. ", + "conclusions": "From analysis of archived radio, X-ray, and gamma-ray data and our own optical spectroscopy and polarimetry, we conclude that 3EG J2006-2321 is a member of the blazar class of AGN. This identification is interesting because it is a reminder that EGRET is capable of detecting blazars with $S_5$ on the order of a fourth of a Jansky. The remaining EGRET unidentified sources most likely to be identified as blazars are in the region $|b|>30^{\\circ}$. Further searches for possible radio and optical counterparts within the error circles of the 30 unidentified sources in this region are encouraged. While the present analysis is sufficient to identify 3EG J2006--2321, little more can be said about the source, and no conclusions are reached regarding beaming and radiation mechanisms of gamma-ray-bright AGN. Complete and simultaneous multiwavelength observations are needed to constrain blazar models. The Gamma-ray Large Area Space Telescope (GLAST), scheduled for launch in 2006, is expected to uncover thousands of gamma-ray blazars and other high-energy sources; however, in order to realize a full return of GLAST science, these sources must be observed not just in gamma rays, but across the electromagnetic spectrum." + }, + "0112/astro-ph0112260_arXiv.txt": { + "abstract": "Nearby dwarf irregular galaxies were searched for compact star clusters using data from the {\\it HST} archives. Eight of the 22 galaxies in our sample were found to host compact clusters of some type. Three of these have populous clusters, with M$_V<-9.5$ at a fiducial age of 10 Myr, and the same three also have super-star clusters, with M$_V<-10.5$ at 10 Myr. Four other dwarf galaxies, two of which contain populous and super-star clusters, are also considered using data in the literature. The results suggest that galaxies fainter than M$_B=-16$ or with star formation rates less than 0.003 M\\solar\\ yr$^{-1}$ kpc$^{-2}$ do not form populous or super-star clusters, and that even the brighter and more active dwarfs rarely form them. Yet when they do form, the associated star formation activity is very high, with numerous compact clusters of similar age in the same complex and evidence for a galaxy-wide perturbation as the trigger. This tendency to concentrate star formation in localized regions of high column density is consistent with previous suggestions that self-gravity must be strong and the pressure must be high to allow a cool phase of gas to exist in equilibrium. Statistical considerations emphasize the peculiarity of super star clusters in dwarf galaxies, which are too small to sample the cluster mass function to that extreme. We suggest that triggered large-scale flows and ambient gravitational instabilities in the absence of shear make the clouds that form super-star clusters in small galaxies. This is unlike the case in spiral galaxies where density wave flows and scale-free compression from turbulence seem to dominate. Further comparisons with spiral galaxies give insight into Larsen \\& Richtler's relation between the star formation rate per unit area and the fraction of young stars in massive dense clusters. We suggest that this relation is the result of a physical connection between maximum cluster mass, interstellar pressure, interstellar column density, and star formation rate, combined with a size-of-sample effect. ", + "introduction": "The properties of super-star clusters set them apart. They are extremely compact and luminous clusters, with full-width at half maximum (FWHM) of less than 15 pc and M$_V$ brighter than about $-$10 (see, for example, van den Bergh 1971; Arp \\& Sandage 1985; Melnick, Moles, \\& Terlevich 1985; Holtzman \\et\\ 1992; Whitmore \\et\\ 1993). If they are bound and contain low mass stars, they resemble what we would expect a globular cluster to be like when it was young. R136, for example, the star cluster in the 30 Doradus \\HII\\ complex in the LMC, is the nearest example of a small super-star cluster. It has a half-light radius \\rhalf\\ of 1.7 pc, an M$_V$ at an age of 2 Myr of $-$11, and a mass of $6\\times10^4$ M\\solar\\ if the mass spectrum continues down to 0.1 M\\solar\\ stars (Hunter \\et\\ 1995). Therefore, the super-star clusters are an extreme mode of star formation. In spite of the fact that the Milky Way has not been able to form a cluster as compact and luminous as a globular cluster for some 10 Gyr (although there may be one forming now---Kn\\\"odlseder 2000), young and old super-star clusters have been found in tiny irregular galaxies. Six super-star clusters are known in five nearby dwarf irregular galaxies and are inferred to be present, though still embedded, in 4 others. NGC 1569, for example, with an integrated M$_V$ of only $-18$, nevertheless, hosts two young super-star clusters with individual M$_V$ of $-$14 and $-$13 (Arp \\& Sandage 1985; O'Connell, Gallagher, \\& Hunter 1994). WLM, with a galactic M$_V$ of only $-$14, contains a bonafide globular cluster that is 14.8 Gyr old (Hodge \\et\\ 1999). So, what are the conditions that allow small irregular galaxies to form super-star clusters when the Milky Way cannot? Do super-star clusters require the same extraordinary circumstances to develop as globular clusters did, and what are those conditions? Populous clusters, less extreme cousins to the super-star clusters, have also been found in large numbers in the Magellanic Clouds (see, for example, Gascoigne \\& Kron 1952; Hodge 1960, 1961; Bica \\et\\ 1996). These objects are less luminous and less massive than the super-star clusters but more extreme in luminosity and compactness than open clusters. A young populous cluster is exemplified by NGC 1818 in the LMC. At an age of about 20 Myr, it has an M$_V$ of $-$9.3, an \\rhalf\\ of 3.2 pc, and a mass of 3$\\times10^4$ M\\solar\\ (Hunter \\et\\ 1997). Besides the Magellanic Clouds, 14 populous clusters have been identified in the irregular NGC 1569 (Hunter \\et\\ 2000), but none have been found in galaxies like IC 1613 (Hodge 1978) and NGC 6822 (Hodge 1977). (The star clusters identified in IC 1613 and NGC 6822 by Hodge are open clusters although one in NGC 6822 has a globular-like metallicity [Cohen \\& Blakeslee 1998]). So, what conditions are necessary for the formation of populous clusters? The high resolution of Hubble Space Telescope {\\it (HST)} is essential to surveying even nearby galaxies for compact clusters. Since a fundamental characteristic of these clusters is their compactness, it is easy to mistake them for a single star with inadequate resolution. The closest super-star cluster, R136 in the LMC, has a half-light diameter of about 13\\arcsec. At a distance of 0.7 Mpc, the lower end of the range studied here, R136 would have a radius of 1\\arcsec, which pushes the limit of typical ground-based telescopes. {\\it HST}'s factor of 10 higher resolution allows us to detect and measure clusters up to 10 times more distant. The problem with resolution is exemplified by the disagreements based on ground-based data over whether the core of R136 was a super massive 2000 M\\solar\\ star or a cluster of very massive stars (Cassinelli, Mathis, \\& Savage 1981) and whether the two super-star clusters in NGC 1569 were really star clusters or bright single stars in that galaxy (Arp \\& Sandage 1985). In both cases the objects are clusters of stars rather than single stars, but studies using ground-based data could not definitely determine this. Due to these inadequacies, a survey for compact clusters beyond the Magellanic Clouds benefits greatly from {\\it HST}. We have used images from the {\\it HST} archives of dwarf irregular galaxies within 7 Mpc to study cluster-producing environments. Unfortunately, since we are using archival data, we did not have a choice in the pointing of the telescope, the filters used, or the exposure times, and the characteristics of what was available placed limits on our search. Promising galaxies that were imaged for a teasing 0.5 s had to be passed over, as did those with only one filter. Only nearby dwarfs galaxies that could be reasonably searched were included in the survey. ", + "conclusions": "\\subsection{General Trends} Galaxies with high star formation rates and high luminosities should have a high abundance of clusters. The top panels of Figure \\ref{figsfr} show the number density of populous and super-star clusters, ND$_{-9.5}$, versus the galactic M$_B$ and the integrated star formation rate normalized to the area of the galaxy SFR$_{25}$. The SFR$_{25}$ is the star formation rate determined from the \\ha\\ emission (Hunter \\& Elmegreen 2002) and normalized to the D$_{25}$ area using the formula of Hunter \\& Gallagher (1986). The open circles are the new galaxies studied here, from Table 1 and the top of Table 2; the triangles are the smallest galaxies in the study by Larsen \\& Richtler (2000; i.e., NGC 1569 and NGC 1705, which are also in our sample); the plus symbols are three of the dwarf irregulars in the literature from the bottom part of Table 2 (NGC 4449, IC 10, and VIIZw 403); and the cross is the WLM galaxy. For galaxies with a range in numbers reflecting our uncertainties in age estimates, we have conservatively chosen the lower end of the range, thereby excluding clusters that would fall into our bin only if an older age were correct. There is a vague correlation between ND$_{-9.5}$ and the galaxy brightness or integrated star formation activity. Most of the galaxies that contain populous clusters are brighter than an M$_B$ of $-16$ and have a $\\log$ SFR$_{25}$ greater than $-2.5$. The galaxy WLM differs because it has a very low current star formation rate ($\\log$ SFR$_{25}$ of $-3.8$) and an M$_B$ of only $-14$, yet it contains am old globular cluster. Presumably the star formation rate in WLM was higher when the globular cluster formed. The cutoffs suggested by Figure \\ref{figsfr} are not the only conditions for populous cluster formation because there are galaxies with values above both cutoffs that have not formed any. In particular, the starburst galaxy IC 10 (Massey \\& Johnson 1998) listed in Table \\ref{tabcld} and the Blue Compact Dwarfs (BCD) VIIZw403 (Lynds \\et\\ 1998) and IZw18 (Hunter \\& Thronson 1995) do not contain populous clusters. IC 10 and VIIZw403 have log SFR$_{25}$ of $-1.5$ and $-2.2$, respectively (Hunter \\& Elmegreen 2002), both above the lower limit for cluster formation, but their mode of star formation has been that of a scaled-up OB association (Hunter \\& Thronson 1995, Lynds \\et\\ 1998, Hunter 2001). \\subsection{The Larsen \\& Richtler study of populous clusters} Larsen \\& Richtler (2000) examined the properties of what they call ``young massive clusters'' in 21 nearby spiral and irregular galaxies supplemented with data from the literature for another 10 objects. Their detection criteria (M$_V=-8.5$ for blue objects and M$_V=-9.5$ for red objects) include populous clusters as well as super-star clusters, and, since the super-star clusters are rare, their sample is dominated by populous clusters. They find that the formation of populous clusters is a normal and on-going process in galaxies with adequate star formation activity. Larsen \\& Richtler (2000) developed a parameter $T_L$(U) to represent the current cluster formation activity in a galaxy; $T_L$(U) is 100 times the ratio of the combined U-band luminosities of the clusters to the U-band luminosity of the whole galaxy. For late-type spiral galaxies, $T_L(U)$ correlates with $M_B$ and SFR. The bottom panels of Figure \\ref{figsfr} show these correlations. The small dots are the regular or starburst spiral galaxies from Larsen \\& Richtler. The smallest two galaxies in their survey (NGC 1705 and NGC 1569) are plotted again as triangles based on data from Tables 1 and 2; they are considered separately because they have about the same diameters as the dwarf galaxies considered here. The largest galaxies (NGC 7252, NGC 3921, NGC 3256, and NGC 1275) are plotted as squares. The open circles and plus symbols are from our dwarf galaxy data, as in the top diagrams. To keep all of the data points on the plot, we let $\\log T_L$(U)$=-2$ if a galaxy has no populous cluster; most of our dwarfs are in this category. The T$_L$(U) for our galaxies were determined from U-band photometry for the galaxies from de Vaucouleurs \\et\\ (1991) and for the clusters from the data presented here. In some cases a U filter was not available for a cluster. In that case we used the available colors and magnitude plus cluster evolutionary models (Leitherer \\et\\ 1999) to predict M$_U$. We do not account for the fact that the {\\it HST} fields do not cover the entire galaxy. Larsen \\& Richtler (2000) noted the trend of $T_L$(U) with M$_B$ and recognized that NGC 1705 and NGC 1569, the two triangles in our plot, have magnitudes that are too faint for their $T_L$(U), as shown also in our figure. Larsen \\& Richtler's relationship between $T_L$(U) and star formation rate is evident in the lower right panel of Figure \\ref{figsfr}. A bi-variate least squares fit to the dots and squares (i.e., all of their galaxies except the two smallest) has a slope of 0.97, as shown by the dashed line. The dwarf irregulars in our survey do not always follow the SFR correlation. Some of the dwarfs with high star formation rates (NGC 1569, NGC 1705, NGC 4449) have populous clusters and follow the trend, but most of the dwarfs do not. WLM currently fits the trend, having $T_L$(U) near zero, but when it was younger the globular cluster would have been much brighter, possibly dominating the U-band emission from the whole galaxy unless the SFR over the entire galaxy was extraordinarily high at the same time. \\subsection{Super-star Clusters} Super-star clusters, with M$_V<-10.5$ at a fiducial age of 10 Myr, are rare in normal spiral galaxies, yet five nearby dwarf irregular galaxies were previously found to contain them (NGC 1569, NGC 1705, NGC 4449, LMC, and WLM), and one additional host was found here: NGC 4214 (Table \\ref{tabssc}). Even more galaxies in Table \\ref{tabcld} could have super-star clusters if an older, but highly uncertain, age is correct (we will not count these here). There are also 4 dwarf galaxies, listed in Table \\ref{tabssc} with a \"?\" mark, that are believed to contain an {\\it embedded} super-star cluster. The properties of these clusters, including their compactness, are inferred from reradiation from the surrounding nebula rather than from direct optical observations of the clusters, so their classification as super-star cluster, rather than a very large OB association, is less certain than if the cluster were visible. The table suggests that galaxies which host a young super-star cluster are all brighter than M$_B=-16$. Most of the super-star clusters are young, with the exception of the globular cluster in WLM. They are also compact, with sizes and masses comparable to globular clusters. Table \\ref{tabssc} lists any crucial status of the galaxy that hosts a super-star cluster. Most of the galaxies show some sign of being affected by an interaction with another galaxy, either currently or in the past. Furthermore, most of the super-star clusters are located near the centers of the galaxies (see also Meurer et al.\\ 1995). A few that are not located near the centers are located at the end of a stellar bar. The centers of interacting galaxies are special places since interactions can drive gas into this region (Noguchi 1988). Similarly, streaming motions around bars can pile the gas up there (Elmegreen \\& Elmegreen 1980). This correlation between the presence of super-star clusters and large-scale galactic disturbances suggests that dwarf irregulars need a special mechanism, like an external perturbation, internal stellar bar, or shell-shell interaction (Chernin, Efremov, \\& Voinovich 1995; Fukui 1999), to concentrate the gas enough to form a massive compact cluster. The exception to all of this is WLM. With an integrated M$_B$ of only $-14$, it nevertheless contains a globular cluster. When the globular cluster was 10 Myr old, it would have had an M$_V$ of order $-14$, rivaling the brightness of the entire galaxy as seen today. Furthermore, WLM does not show any particular sign of having been disrupted by an interaction 15 Gyr ago or of having gone through an extraordinary starburst then, although our ability to probe the star formation histories that far back is limited. The existence of this globular cluster in such a small galaxy and its rarity does suggest that something extraordinary happened to WLM nearly 15 Gyr ago when the galaxy was first forming. Other dwarf irregular galaxies are undergoing starbursts that do not show evidence of compact star formation and super-star clusters. IC 10, for example, is a starburst galaxy (Massey \\& Johnson 1998) that has not produced anything but small clusters and OB associations at least in the portion imaged with {\\it HST} (Hunter 2001). VIIZw403 and IZw18 are BCDs that have only made large OB associations (Hunter \\& Thronson 1995, Lynds \\et\\ 1998). Neither the conditions that produce a starburst nor the starburst itself are sufficient to produce clusters of such an extreme nature. We do find a rather gregarious nature to super-star clusters. The super-star cluster in NGC 1705 has a collection of 10 smaller compact clusters arranged in a semi-circle around it. Six of these have an age that is similar to that of the super-star cluster. There are also 10 compact clusters surrounding the youngest of the two super-star clusters in NGC 1569 (Hunter \\et\\ 2000). Again, at least 5 of these have an age comparable to that of the super-star cluster and another 3 have ages only $\\sim$10 Myr older. There are also OB associations formed over the past 10 Myr arrayed around the super-star cluster R136 in the LMC that is itself 2 Myr old. It appears that these super-star clusters have formed from very large gas complexes that produced compact clusters and OB associations as associated debris. Star formation in such complexes appears to extend over a time-scale of roughly 10 Myr. However, not all such complexes form super-star clusters (Kennicutt \\& Chu 1988). Constellation III in the LMC, for example, is an older complex that formed stars over a scale comparable to the 30 Doradus nebula in which R136 sits. However, only a few populous clusters have formed in Constellation III (Dolphin \\& Hunter 1998). After examining {\\it HST} images of 22 nearby dwarf irregular galaxies, we have found compact star clusters in 8 of them. Among these 8 we found 3 galaxies (NGC 4214, NGC 1705, WLM) that host populous clusters, which we define as clusters with M$_V$ at a fiducial age of 10 Myr brighter than $-$9.5. These same three galaxies host super-star clusters, which we define as clusters with M$_V$ at an age of 10 Myr brighter than $-10.5$. Another 4 galaxies in our survey might contain populous clusters and another 3 could join the list as containing super-star clusters {\\it if} the older age estimate in a range of ages consistent with the cluster's color is correct. However, this seems unlikely. We supplement our survey with our previous census of compact clusters in IC 10, NGC 1569, NGC 4449, and VIIZw403. IC 10, NGC 1569, and NGC 4449 contain compact clusters candidates, but only NGC 1569 and NGC 4449 contain populous clusters, and both also contain super-star clusters. >From these host galaxies, if we exclude WLM, we find a magnitude cutoff of M$_B=-16$ for the formation of populous and super-star clusters; fainter than this and a galaxy does not seem to form these extreme star clusters with a high enough probability to have detected any in our sample. This is consistent with Hodge \\et's (1999) claim that no dwarf irregular galaxy, except WLM, in the Local Group with M$_V$ fainter than $-16$ contains a globular cluster. There is also a minimum normalized star formation rate of 0.003 M\\solar\\ yr$^{-1}$ kpc$^{-2}$ for cluster formation. These limits, however, are not sufficient to mandate the development of clusters; there exist bright enough galaxies with high star formation rates that are barren. The dwarf irregulars also contribute to the trends seen for spiral galaxies by Larsen \\& Richtler (2000): The contribution of populous clusters to the U-band light of a galaxy increases as the integrated galactic M$_B$ and normalized star formation rate increases. However, a few of the dwarf irregulars also contribute to the scatter and, again, there are examples of galaxies with high M$_B$ or star formation rate and zero contribution from populous clusters. We can explain these results in terms of a general framework for understanding the Larsen \\& Richtler (2000) $T_L(U)$ correlations, combined with the small sizes of dwarf galaxies and the observation that most dwarfs with super-star clusters look agitated. The correlations follow partly from the inter-dependencies of pressure, gas column density, and star formation rate in galaxies that are all about the same size and also large enough to sample the cluster mass spectrum out to the maximum cluster mass that is likely to form. Small galaxies cannot generally sample out this far, and will usually fall short of producing a massive compact cluster. Small galaxies with global perturbations can make massive bound clusters, however. They may do this by the same star formation processes that operate in spiral galaxies, in which case the dwarfs with super star clusters are stochastic fluctuations in the dwarf sample. Or, the dwarf may produce super star clusters by a mechanism that is unique to their type, such as a gravitational instability that is rendered more effective by the low rate of shear. Similar processes may occur in the nuclear regions of large galaxies." + }, + "0112/astro-ph0112110_arXiv.txt": { + "abstract": "The infrared (IR) emission from interstellar dust in the Small Magellanic Cloud (SMC) is modelled using a mixture of amorphous silicate and carbonaceous grains, including a population of polycyclic aromatic hydrocarbon (PAH) molecules. (1) It is shown that this dust model is able to reproduce the spectral energy distribution from near-IR to far-IR for the entire SMC Bar region, provided the PAH abundance in the SMC Bar region is very low. (2) The IR spectrum of the \\smcb\\ molecular cloud can also be reproduced by our dust model provided the PAH abundance is increased relative to the overall SMC Bar. The PAHs in \\smcb\\ incorporate $\\sim$3\\% of the SMC C abundance, compared to $<0.4\\%$ in the SMC Bar. (3) The spectrum of \\smcb\\ is best reproduced if the PAH mixture has intrinsic IR band strengths which differ from the band strengths which best fit Milky Way PAH mixtures. The variation in the PAH IR band strengths would imply different PAH mixtures, presumably a consequence of differing metallicity or environmental conditions. Other possibilities such as super-hydrogenation of PAHs and softening of the starlight spectrum are also discussed. ", + "introduction": "} Due to its low metallicity ($\\sim$1/10 of that in the Milky Way [Dufour 1984]) and low dust-to-gas ratio (over 10 times lower than in the Milky Way [Bouchet et al.\\ 1985]), the Small Magellanic Cloud (SMC) is often considered as a prototype for the interstellar matter in high-redshift galaxies at early stages of chemical enrichment. The analogy is strengthened by the similarities of its typical extinction curve -- a nearly linear rise with $\\lambda^{-1}$ ($\\lambda$ is wavelength) and the absence of the 2175\\AA\\ hump (Lequeux et al.\\ 1982; Pr\\'evot et al.\\ 1984; Bouchet et al.\\ 1985; Thompson et al.\\ 1988; Rodrigues et al.\\ 1997; Gordon \\& Clayton 1998) -- to wavelength-dependent extinction inferred for starburst galaxies (Gordon, Calzetti, \\& Witt 1997) and damped Ly$\\alpha$ systems (Pei, Fall, \\& Bechtold 1991). Infrared (IR) emission from dust in the SMC has been seen by the {\\it Infrared Astronomical Satellite} (IRAS) (Schwering \\& Israel 1990; Sauvage, Thuan, \\& Vigroux 1990), the Diffuse Infrared Background Experiment (DIRBE) instrument on the {\\it Cosmic Background Explorer} (COBE) satellite (Stanimirovic et al.\\ 2000), and the {\\it Infrared Space Observatory} (ISO) (Reach et al.\\ 2000). In the present work we use these observations to test the grain model proposed previously to account for the observed extinction curve for dust in the general SMC Bar region (Weingartner \\& Draine 2001a). Of particular interest is the detection by Reach et al.\\ (2000) of the 6.2, 7.7, 8.6, 11.3, and 12.7$\\mu$m emission bands toward the quiescent molecular cloud \\smcb. These bands, generally attributed to polycyclic aromatic hydrocarbons (PAHs) (L\\'eger \\& Puget 1984; Allamandola, Tielens, \\& Barker 1985), had not previously been observed in the SMC. The molecular cloud \\smcb\\ ($[\\alpha,\\delta]_{1950}=[0^{\\rm h}43^{\\rm m}42^{\\rm s}4, -73^{\\rm o}35^{\\prime}10^{\\prime\\prime}]$) has a 11.3$\\mu$m/7.7$\\mu$m emission band ratio which differs significantly from that for Milky Way regions, such as the $\\rho$ Oph reflection nebula (which one might expect SMC B1\\#1 to resemble [see Reach et al.\\ 2000]); the observed band ratio also falls outside the range which appears to be ``allowed'' by the PAH model of Li \\& Draine 2001 (see Figures 16,17 of Draine \\& Li 2001). The observed spectrum for \\smcb\\ therefore presents an important test which may reveal shortcomings of the current PAH emission models. The unusual band ratio might be indicative of the physical conditions in \\smcb, or it might point to systematic differences between the PAH mixtures in the Milky Way and SMC. In this work we extend the interstellar dust emission model developed for the diffuse interstellar medium (ISM) of the Milky Way (Li \\& Draine 2001; hereafter LD01) to the SMC, with emphasis on the mid-IR spectrum of \\smcb. We seek (1) to test the PAH IR emission model (Draine \\& Li 2001; LD01); (2) to infer the SMC PAH properties such as size distributions, charging, hydrogenation, and intrinsic IR band strengths; (3) to infer the physical conditions in the SMC and the \\smcb\\ cloud and the dust properties in these environments. Our approach is outlined in \\S\\ref{sec:dustmod} and \\S\\ref{sec:parameters}. We first model the spectral energy distribution for the SMC as a whole in \\S\\ref{sec:smcsed}. Model results for the \\smcb\\ molecular cloud are presented in \\S\\ref{sec:mwmod}--\\S\\ref{sec:ir_band_strengths}. In \\S\\ref{sec:extinction} we discuss the 2175\\AA\\ extinction hump predicted from the present PAH model. Our principal conclusions are summarized in \\S\\ref{sec:sum}. ", + "conclusions": "} In contrast to the Milky Way Galaxy, the SMC is characterized by a lower metallicity, a lower dust-to-gas ratio, and an absence of the 2175\\AA\\ hump on most sightlines with measured UV extinction (see \\S\\ref{sec:intro}). Therefore, the SMC provides an ideal laboratory to study dust in an environment quite different from the Milky Way. We have modelled the IR spectral energy distribution for the SMC as a whole (averaged over a 6.25 deg$^2$ area) (\\S\\ref{sec:smcsed}). We have also modelled the IR emission spectrum of the quiescent molecular cloud \\smcb\\ (\\S\\ref{sec:smcb1sed}). Our principal results are: \\begin{enumerate} \\item The dust IR emission from a 6.25 deg$^2$ region including the SMC Bar and Eastern Wing is well reproduced by a model with $\\dN/d\\Urad\\propto \\Urad^{-1.8}$, $10^{-1.0} \\le \\Urad \\le 10^{2.75}$, and $N_{\\rm H}^{\\rm tot}\\approx 5.4\\times 10^{21}\\cm^{-2}$ using the SMC dust model of Weingartner \\& Draine (2001a). This dust has a very low PAH abundance, with C/H $\\ltsim$ 0.2ppm ($\\ltsim 0.4\\%$ of the SMC interstellar C abundance) incorporated into PAHs. The dust-weighted mean radiation intensity is $\\langle U\\rangle \\approx 2$, but the starlight intensity must have a broad distribution ranging from $\\Umin\\ltsim 0.3$ to $\\Umax\\gtsim 50$ (\\S\\ref{sec:smcsed}). \\item The IR emission from the molecular cloud \\smcb\\ can be reproduced by a dust model consisting of silicates, graphite, and PAHs with $N_{\\rm H}^{\\rm tot}\\approx 4.2\\times 10^{22}\\cm^{-2}$ illuminated by starlight with a power law distribution $\\dN/d\\Urad\\propto \\Urad^{-2}$, $10^{-0.5} \\le \\Urad \\le 10^{2.5}$ (\\S\\ref{sec:smcb1sed}). The required PAH abundance in \\smcb\\ is 1.5$\\ppm$, or $\\sim$3\\% of the SMC interstellar C abundance. \\item The observed 11.3$\\um$/7.7$\\um$ and 6.2$\\um$/7.7$\\um$ band ratios for \\smcb\\ fall outside the ``allowed'' region predicted for the Milky Way PAH model (see Figures 16, 17 of Draine \\& Li 2001). In \\S\\ref{sec:mwmod}, \\S\\ref{sec:h2c}, and \\S\\ref{sec:reddenedisrf}, we examine the possibility that this could be due to (1) an enhanced neutral fraction, (2) an enhanced H/C ratio for the PAHs in \\smcb, or (3) reddening of the radiation field to which the PAHs in \\smcb\\ are exposed, but none of these scenarios successfully accounts for the observed 5--15$\\micron$ spectrum. We instead conclude that the intrinsic IR band strengths for the PAH mixture in \\smcb\\ differ from the band strengths adopted for Milky Way PAH mixtures (\\S\\ref{sec:ir_band_strengths}). This indicates that the PAH mixture in \\smcb\\ differs in detail from the PAH mixtures present in a number of Milky Way regions. The band strengths suggested by the \\smcb\\ spectrum are close to ``laboratory'' band strengths. \\item The PAH abundance in \\smcb\\ exceeds that in the SMC Bar by a factor $\\gtsim 8$ (C/H=1.5 ppm vs $\\ltsim$ 0.2ppm). While the average SMC extinction curve has no detectable 2175\\AA\\ hump, we predict a 2175\\AA\\ extinction excess of $\\simgt 0.5\\magni$ for the dust in \\smcb\\ (\\S\\ref{sec:extinction}). \\end{enumerate} Two types of measurements would be of great value to test our dust model: (1) UV extinction measurements for sightlines through or near \\smcb\\ to test for the predicted 2175\\AA\\ feature; (2) $3-15\\micron$ spectrophotometry of other regions in the SMC where UV extinction measurements place a strong upper limit on the 2175\\AA\\ feature -- the present model predicts that PAH emission should not be seen from such regions." + }, + "0112/astro-ph0112326_arXiv.txt": { + "abstract": "We have used observations of Cygnus X-1 from the Compton Gamma-Ray Observatory (\\gro\\/) and BeppoSAX to study the variation in the MeV $\\gamma$-ray emission between the hard and soft spectral states, using spectra that cover the energy range 20 keV up to 10 MeV. These data provide evidence for significant spectral variability at energies above 1 MeV. In particular, whereas the hard X-ray flux {\\it decreases} during the soft state, the flux at energies above 1 MeV {\\it increases}, resulting in a significantly harder $\\gamma$-ray spectrum at energies above 1 MeV. This behavior is consistent with the general picture of galactic black hole candidates having two distinct spectral forms at soft $\\gamma$-ray energies. These data extend this picture, for the first time, to energies above 1 MeV. We have used two different hybrid thermal/non-thermal Comptonization models to fit broad band spectral data obtained in both the hard and soft spectral states. These fits provide a quantitative estimate of the electron distribution and allow us to probe the physical changes that take place during transitions between the low and high X-ray states. We find that there is a significant increase (by a factor of $\\sim 4$) in the bolometric luminosity as the source moves from the hard state to the soft state. Furthermore, the presence of a non-thermal tail in the Comptonizing electron distribution provides significant constraints on the magnetic field in the source region. ", + "introduction": "High energy emission from galactic black hole candidates (GBHCs) is characterized by variability on time scales ranging from msec to months. In the case of Cygnus X-1, it has long been recognized that, on time scales of several weeks, the soft X-ray emission ($\\sim 10$ keV) generally varies between two discrete levels \\citep[e.g.,][]{priedhorsky1983,ling1983,liang1983}. The source seems to spend most ($\\sim 90\\%$) of its time in the so-called {\\em low X-ray state}, characterized by a relatively low flux of soft X-rays and a relatively high flux of hard X-rays ($\\sim100$ keV). This state is sometimes referred to as the {\\em hard state}, based on the nature of its soft X-ray spectrum. On occasion, it moves into the so-called {\\em high X-ray state}, characterized by a relatively high soft X-ray flux and a relatively low hard X-ray flux. This state is sometimes referred to as the {\\em soft state}, based on the nature of its soft X-ray spectrum. There are, however, some exceptions to this general behavior. For example, HEAO-3 observed, in 1979, a relatively low hard X-ray flux coexisting with a low level of soft X-ray flux \\citep{ling1983,ling1987}. \\citet{ubertini1991} observed a similar behavior in 1987. Observations by the BATSE, OSSE, COMPTEL and EGRET instruments on the Compton Gamma-Ray Observatory (\\gro\\/), coupled with observations by other high-energy experiments (e.g., SIGMA, ASCA and RXTE) have provided a wealth of new information regarding the emission properties of galactic black hole candidates. One important aspect of these high energy radiations is spectral variability, observations of which can provide constraints on models which seek to describe the global emission processes. Based on observations by OSSE of seven transient galactic black hole candidates at soft $\\gamma$-ray energies (i.e., below 1 MeV), two $\\gamma$-ray spectral shapes have been identified that appear to be well-correlated with the soft X-ray state \\citep{grove1997,grove1998,grove1999}. In particular, these observations define a {\\it breaking} $\\gamma$-ray spectrum that corresponds to the hard (low) X-ray state and a {\\it power-law} $\\gamma$-ray spectrum that corresponds to the soft (high) X-ray state. A thorough understanding of the nature of these systems requires modeling that cannot only explain the individual spectra, but can also explain the transitions between the various spectral states \\citep[e.g.,][]{grove1998,liang1998,poutanen1998a}. In recent years, a general theoretical picture of the accretion flow in Cygnus X-1 has emerged which appears to provide a reasonable explanation of the spectral data in both the low and high X-ray states. This model includes an inner optically-thin, geometrically-thick advection-dominated accretion flow (ADAF) surrounded by an outer, geometrically-thin, optically-thick accretion disk \\citep{esin1998}. The outer disk is characterized by a blackbody spectrum. ADAF flows \\citep[e.g.,][]{narayan1996} are characterized by their relatively low radiative efficiencies and by a two-temperature structure, with the ions nearly virial at $T_i \\sim10^{12}$ K and the electrons at $T_e \\sim 10^9$ K. The high temperature of the ADAF leads to an extended, quasi-spherical geometry. Hot optically-thin ADAFs exist only below a certain critical accretion rate. The transition radius between the ADAF and the thin disk therefore depends on the accretion rate. At higher accretion rates, where it is more difficult to support the ADAF, the transition radius moves to smaller radii, closer to the black hole. The ADAF region is largely responsible for the hard X-ray flux ($\\sim20-100$ keV), while the outer thin disk is generally responsible for the soft X-ray flux ($\\sim2-10$ keV). In the context of this general model, the spectral state of Cygnus X-1 depends on the accretion rate. At low accretion rates, the inner ADAF extends out to a transition radius of $\\sim100$ Schwarzschild radii \\citep{esin1998}. In this configuration, the ADAF region makes a significant contribution to the hard X-ray flux. At higher accretion rates, it becomes more difficult to support the ADAF. The ADAF region therefore shrinks, and the transition radius moves inward, although there may exist a low-density ADAF corona surrounding the thin disk \\citep{narayan1998}. The level of hard X-ray flux decreases due to the smaller volume of the ADAF region, while the level of soft X-ray flux increases due to the larger size of the thin disk region. In this scenario, the hard state corresponds to a relative low accretion rate, with the spectrum dominated by the ADAF region, and the soft state corresponds to a relatively high accretion rate, with the spectrum dominated by the blackbody of the outer thin disk region. Small changes in the accretion rate (on the order of 10--15\\%) may be sufficient to trigger a transition between the hard and soft states \\citep{esin1998}. The ADAF model described above provides a consistent framework for understanding the essential dynamics and spectra of black hole accretion flows. In the context of this framework, however, simple thermal Comptonization models appear unable to account for all of the spectral features, especially the hard power-law tail that is seen at energies above $\\sim600$ keV \\citep{gierlinski1999}. \\citet{poutanencoppi1998} used a geometry similar to that described above \\citep{poutanen1997} and assumed some (unspecified) source of non-thermal electrons that remains constant during the spectral state transitions. This suggests that the non-thermal component may play a more significant role, especially at higher energies, during the high X-ray state, where the ADAF contribution is suppressed. Hybrid thermal/non-thermal plasmas have often been successfully used to model the observed data \\citep[e.g.,][]{gierlinski1999,poutanencoppi1998}. Based on the assumption that the spectrum results from inverse Compton scattering of a thermal photon spectrum by energetic electrons, the underlying electron population could be described as a combination of a thermal Maxwellian and a power-law tail extending to higher energies. The presence of a non-thermal component is often assumed a priori, without any specific model to explain the origin, although the existence of such distributions is clearly established in the case of solar flares \\citep[e.g.,][]{coppi1999} and it is therefore natural to expect that similar distributions exist elsewhere in the universe \\citep[e.g.,][]{crider1997,gierlinski1997,poutanen1996,poutanen1998a,poutanencoppi1998,coppi1999}. Others have considered physical mechanisms by which non-thermal electron distributions might be developed. For example, both stochastic particle acceleration \\citep{dermer1996,li1996} and MHD turbulence \\citep{li1997} have been proposed as mechanisms for directly accelerating the electrons. The ion population might also contribute to the non-thermal electron distribution in the case where a two-temperature plasma develops \\citep[e.g.,][]{dahlbacka1974,shapiro1976,chakrabarti1995}. With ion population temperatures approaching $kT_i \\sim 10^{12}$ K, $\\pi^o$ production from proton-proton interactions may take place \\citep[e.g.,][]{eilik1980,eilik1983,mahadevan1997}. The $\\pi^o$ component may then lead, via photon-photon interactions between the $\\pi^o$-decay photons and the X-ray photons, to production of energetic (nonthermal) $e^{+}-e^{-}$ pairs. \\citet{jourdain1994} used this concept to fit the hard X-ray tails of not only Cygnus X-1, but also GRO J0422+32 and GX 339-4, as measured by both SIGMA and OSSE. While retaining a standard thermal Comptonization spectrum \\citep{sunyaev1980} to explain the emission at energies below 200 keV, they used $\\pi^o$ production to generate the nonthermal pairs needed to fit the spectrum at energies above $\\sim 200$ keV. The power-law spectra seen in the high X-ray state have also been modeled as resulting from bulk-motion Comptonization \\citep[e.g.,][]{ebisawa1996b,titarchuk1997,laurent1999}. In this model, the flow becomes quasi-spherical within the inner-most stable orbit. The nearly relativistic flow of the free-falling electrons gives rise to the Comptonization of ambient photons. This model predicts power-law spectra, with a slope that depends on the mass acretion rate. The difficulty with this model is that it predicts spectral sharp cutoffs below 500 keV, a result that is clearly inconsistent with the observed spectra. Although we cannot rule out bulk motion Comptonization as a contributor to the spectrum at lower energies, it is clearly not capable of accounting for the high energy emission. Improvements in the theoretical modeling of spectral state transitions can be expected to arise from improved observations at energies above 600 keV. It will be important to understand how this part of the spectrum, most likely dominated by non-thermal emission, changes during the spectral transition. Of particular interest will be observations that can discern a clear cutoff in the spectra at high energies. The precise energy of the cutoff is a function of the compactness of the source region, since it is influenced by $\\gamma-\\gamma$ opacity. A measure of the cutoff energy, possibly coupled with measurements of the 511 keV $e^{\\pm}$ annihilation line, will help constrain the compactness of the region responsible for the emission and determine the extent to which $e^{\\pm}$ pairs may play a role in the emission region \\citep{poutanen1998a}. Using hard state data collected during the first three years of the \\gro\\/ mission, \\citet{mcconnell2000a} compiled a broad-band hard state spectrum of Cygnus X-1 using contemporaneous data from all four instruments on \\gro\\/ (BATSE, OSSE, COMPTEL and EGRET). Unlike previous broad-band studies, these data provided a measurement of the spectrum at energies above 1 MeV. The resulting spectrum showed evidence for significant levels of non-thermal emission at energies out to 5 MeV. The spectral shape, although consistent with the so-called breaking spectral state \\citep{grove1997,grove1998} of the $\\gamma$-ray emission, was clearly not consistent with standard Comptonization models. The hybrid thermal/non-thermal model of \\citet{poutanen1996} was used to fit the hard state data, with fits that indicated a thermal electron population with a temperature of $\\sim90$ keV and a high energy power-law electron component with a spectral index of $\\sim4.5$. In May of 1996, a transition of Cyg X-1 into a soft state was observed by RXTE, beginning on May 10 \\citep{Cui1997}. The 2--12 keV flux reached a level of 2 Crab on May 19, four times higher than its normal value. Meanwhile, at hard X-ray energies (20-200 keV), BATSE measured a significant {\\it decrease} in flux \\citep{zhang1997}. Motivated by these dramatic changes, a target-of-opportunity (ToO) for \\gro\\/, with observations by OSSE and COMPTEL began on June 14 (\\gro\\/ viewing period 522.5). Here we report on the results from an analysis of the \\gro\\/ data from this ToO observation, incorporating the high energy results from COMPTEL. This includes a comparison with results obtained from an updated analysis of \\gro\\/ soft state data, making use of the same data studied previously by \\citet{mcconnell2000a}. In $\\S$ 2 we describe the \\gro\\/ observations of Cygnus X-1 in its hard state. The data analysis is described in $\\S$ 3, followed by a discussion of those results in $\\S$ 4. ", + "conclusions": "The COMPTEL data alone can be used to draw some important conclusions regarding the MeV variability of Cyg X-1. Most importantly, the flux measured by COMPTEL at energies above 1 MeV was observed to be {\\em higher} (by a factor of 2.5) during the soft state (in May of 1996) than it was during the hard state (as averged over several \\gro\\/ observations). This is in contrast to the {\\em lower} flux level observed at hard X-ray energies (i.e., near 100 keV) during the soft state. The lack of any detectable emission by COMPTEL below 1 MeV (i.e., in the 750 keV to 1 MeV energy band) further suggests a hardening of the $\\gamma$-ray spectrum during the soft state. Inclusion of the BATSE and OSSE spectra adds considerably more information regarding the spectral variability. Whereas the low-state \\gro\\/ spectrum shows the breaking type spectrum that is typical of most high-energy observations of Cyg X-1 \\citep[e.g.,][]{mcconnell2000a}, the high-state \\gro\\/ spectrum shows the power-law type spectrum that is characteristic of black hole candidates in their high X-ray state. Our analysis of the soft state data from BATSE, OSSE and COMPTEL shows that the spectrum at these energies can be described by a single power-law with a best-fit photon spectral index of $\\Gamma = 2.58 \\pm 0.03$. A similar spectrum had already been reported for this same time period (VP 522.5) based on independent studies with data from both BATSE \\citep{zhang1997} and OSSE \\citep{gierlinski1997,gierlinski1999}. A detailed study of the broadband soft state spectrum, based on data from ASCA, RXTE and \\gro\\//OSSE, was reported by \\citet{gierlinski1999}, but they did not include the higher energy COMPTEL data. The inclusion of the COMPTEL data in the high state spectrum provides evidence, for the first time, of a continuous power-law (with a photon spectral index of 2.6) extending beyond 1 MeV, up to $\\sim10$ MeV. No clear evidence for a cutoff in the power-law spectrum can be discerned from these data. A power-law spectrum had also been observed by both OSSE and BATSE during the high X-ray state of February, 1994 \\citep[\\gro\\/ VP 318.1; ][]{phlips1996,ling1997}. These earlier data correspond to the low level of hard X-ray flux near TJD 9400 in Figure 1. The spectrum observed during the 1994 high state showed a similar photon spectral index ($\\Gamma = 2.72$ vs. $\\Gamma = 2.57$ for the 1996 high state spectrum), but the overall intensity of the power law was considerably lower \\citep{gierlinski1999}. Near 1 MeV, for example, the spectral amplitude was about 3 times lower in 1994 than it was in 1996. This explains why Cygnus X-1 was not observed by COMPTEL during the 1994 high state. The extrapolation of the lower-intensity power-law fell below the sensitivity limit of COMPTEL. On the other hand, the intensity observed in 1996 was sufficiently high to allow for a measurement of the spectrum by COMPTEL. We have used two different hybrid thermal/non-thermal Comptonization models ({\\tt compps} and {\\tt eqpair}) to fit broad band spectral data obtained in both the hard and soft spectral states. For the hard state analysis, we used data from \\gro\\/ covering 20 keV up to 10 MeV. For the soft state analysis, we augmented the \\gro\\/ data with lower energy data from BeppoSAX to provide improved constraints on the spectrum at energies down to 0.5 keV. These fits provide a quantitative estimate of the electron distribution and allow us to probe the physical changes that take place during transitions between the low and high X-ray states. Hybrid Comptonization models have also been used to model the spectra of other black hole binaries in their soft state, such as GRS 1915+105 \\citep{aaz01}. The high energy spectrum of Cygnus X-1 cannot be described by the bulk-motion Comptonization model alone, which predicts a sharp cutoff above $\\sim 100$ keV (Laurent \\& Titarchuk 1999). The hybrid comptonization models provide an adequate fit to the data without requiring any contribution from bulk-motion Comptonization. Furthermore, the bulk-motion Comptonization power-law for $L\\sim 0.04 L_{\\rm E}$ corresponding to the soft state of Cygnus X-1 (see below), was found by Laurent \\& Titarchuk (1999) to be very soft, with $\\Gamma\\simeq 3.5$ at , i.e., much softer than the observed $\\Gamma\\simeq 2.5$ \\citep{gierlinski1999,frontera2001}. (See also the discussion in Zdziarski 2000.) Note that the XSPEC model of bulk-motion Comptonization, {\\tt bmc} (Shrader \\& Titarchuk 1998), does not include any high-energy cutoff and thus cannot be applied to our data (or any data extending to $\\ga 100$ keV). Figure \\ref{f:compare} shows a comparison of the spectra in the two states. For the hard state, we also show a typical spectrum at energies $\\la 25$ keV \\citep[\\sax\\/ data from][]{disalvo2001}. We see that the two broad-band spectra cross each other at $\\sim 10$ keV and $\\sim 1$ MeV. The dashed curve shows the model obtained by fitting the hard-state data from \\gro\\/ only (\\S \\ref{s:hard}), and assuming $\\nh=6\\times 10^{21}$ cm$^{-2}$. We see that this model predicts the low-energy \\sax\\/ data relatively well, underestimating somewhat the observed spectrum only at $\\la 10$ keV due to the presence of a pronounced soft X-ray excess present in the hard state \\citep{ebisawa1996a,frontera2001,disalvo2001}, which is neglected in our model fitted to the data at $\\geq 20$ keV. The bolometric flux or luminosity ratio between the soft state in June 1996 and the average for the hard state is $\\sim 4$. This value is much more than the rough estimate of $\\sim 1.5-1.7$ based on the ASM and BATSE occultation results \\citep{zhang1997}, but is consistent with the results of \\citet{frontera2001}, based on studies with BeppoSAX. Such a large value makes models of the state transition based on a change of accretion rate plausible. Given the larger luminosity in the soft state, the characteristic dimension of the hot plasma in the soft state based on the compactness fit is similar to that in the hard state, $\\sim 10^2 GM/c^2$. These data tend to support the general picture that the transition between the hard and soft states results from a change in the disk transition radius between a hot inner corona (ADAF) and a cooler outer thin disk \\citep[e.g.,][]{esin1998,narayan1998,poutanen1998a,poutanen1998b,poutanencoppi1998}. In the hard state, this transition radius is relatively far from the black hole (at $\\sim 100$ Schwarzschild radii). The spectrum is dominated by Comptonization off the thermal electrons in the hot inner corona. Radio emission is also more pronounced in this state \\citep{fender2001}, with evidence for a radio emitting relativistic jet \\citep{stirling2001}. As the transition radius moves inward, perhaps due to an increase in the accretion rate, the optically thick cool disk intercepts a larger fraction of the energy. The thermal energy dissipation in the corona is reduced considerably and the blackbody disk component (the principal component at soft X-ray energies) becomes more pronounced. Although our data tend to support the above picture, we have not attempted to model the geometry in detail, since the precise geometrical configuration of the emitting region is largely unknown. Furthermore, our new data cover the energy range near 1 MeV where geometry effects are difficult to study. One of the primary goals of the present paper is to determine the electron distribution of the radiating plasma. Our assumption of a spherical source geometry provides the necessary physics that is required to extract information on the electron spectrum. We have further presumed that the thermal and nonthermal electrons are in the same physical region. This assumption is based, in part, on the observations that show a negative correlation between the thermal and nonthermal components. This need not be the case in reality, however, but the present data cannot be used to determine the extent to which the two populations are co-located. A more detailed discussion of geometrical effects in the context of the \\texttt{eqpair} model, including Compton reflection and energy balance, can be found in \\citet{gierlinski1999}. The shape of the electron distribution and its high energy tail can best be determined by measurements that extend into the MeV energy region. The high-energy cutoff is related to the compactness of the source region, since it depends, in part, on the influence of $\\gamma-\\gamma$ pair production. If $\\gamma-\\gamma$ pair production is an important source of opacity, this would imply the presence of a significant level of $e^{\\pm}$ pairs in the source region. In this way, a measure of the high-energy cutoff can help determine the nature of the emitting plasma ($e-p$ or $e^{\\pm}$). Although a measure of $e^{\\pm}$ annihilation radiation can also serve as a diagnostic of a pair plasma, it is likely that any annihilation radiation that may be present would be considerably broadened (and perhaps blue-shifted), and hence may not be readily observable. Measurements to date with HEAO-3 \\citep{ling1989} and with OSSE \\citep{phlips1996} provide only upper limits, or, at best, a marginal ($1.9\\sigma$) detection \\citep{ling1989} to the level of $e^{\\pm}$ annihilation radiation. This further underscores the need to define the high-energy cutoff as perhaps the best means for constraining the source compactness and the nature of the emitting region. If INTEGRAL, with its improved line sensitivity, succeeds in measuring an annihilation feature, then constraints on the high-energy cutoff will be even more valuable. The presence of a non-thermal tail in the electron distribution can also provide constraints on the strength of the magnetic field in the source region. As pointed out by Wardzi\\'nski \\& Zdziarski (2001), the presence of even a weak nonthermal electron tail increases strongly the emissivity of the cyclo-synchrotron process with respect to the pure thermal case. If the Compton-scattering electrons in Cygnus X-1 were purely thermal, that process appears in general to be too inefficient to provide all of seed photons for the Comptonization under simple assumptions of equipartition (Wardzi\\'nski \\& Zdziarski 2000). Since we do see a blackbody component at low energies \\citep{ebisawa1996a,disalvo2001}, this inefficiency is consistent with the seed photons for Comptonization provided by the blackbody rather than by the cyclo-synchrotron photons. On the other hand, the tail parameters obtained by \\citet{mcconnell2000a} yielded such a copious supply of cyclo-synchrotron seed photons that the corresponding luminosity would become $\\sim 10^2$ times that observed (Wardzi\\'nski \\& Zdziarski 2001). This conclusion is confirmed for the tail parameters fitted here (G. Wardzi\\'nski, private communication). Thus, either the magnetic field in Cygnus X-1 is substantially below equipartition (at least an order of magnitude) or the observed photon tail has a different origin than that due to a high energy electrons. In either case, this has important implications for models of the accretion flow in Cygnus X-1. These studies also have implications that go beyond that of studying individual black hole sources. Given the close spectral similarity between black-hole binaries in the hard state and Seyferts (e.g.,\\ Zdziarski 2000), it is possible that similar tails are present in the spectra of the latter objects. \\citet{stecker1999} have suggested that the hard tail emission seen in sources like Cygnus X-1 might account for an important component of the cosmic diffuse background radiation in the 200 keV -- 3 MeV energy band \\citep[see also][]{stecker2001}. Note, however, that the tail of Cygnus X-1 above 1 MeV contains relatively little flux, 1.3\\% of the bolometric (model) flux, for the fit with {\\tt eqpair}. If a similar value is characteristic of Seyferts, the combined emission from their high energy tails may be too weak to account for the observed extragalactic MeV background, perhaps arguing against the proposal by \\citet{stecker1999}. The next major satellite for this energy range, INTEGRAL, is expected to have only slightly better continuum sensitivity than COMPTEL at energies near 1 MeV with both its IBIS and SPI experiments \\citep{schoenfelder2001}. Furthermore, the much narrower FoV of the INTEGRAL instruments ($\\sim15\\degr$) will mean that there will likely be only a limited number of observations of Cygnus X-1. This is in stark contrast to the COMPTEL situation, in which the large FoV of COMPTEL ($\\sim60\\degr$) resulted in many weeks of exposure, most of which were obtained during the low X-ray state. Given the large low-state exposure of \\gro\\/, it is quite likely that INTEGRAL may not be able to offer any significant improvement in our knowledge of the hard state continuum spectrum at MeV energies. The \\gro\\/ data may therefore provide the best view of the hard state MeV continuum for many years to come. However, COMPTEL is very limited in the data that it collected for the {\\em soft state} spectrum. Additional soft state observations with INTEGRAL could therefore prove valuable. An important goal would be to search for a cutoff in the energy spectrum. Pinning down the energy of this cutoff would be a very important next step in our understanding of the high energy spectrum of Cygnus X-1. In this regard, INTEGRAL may be an extremely useful tool for collecting additional soft state spectral data, providing that suitable target-of-opportunity observations can be acquired." + }, + "0112/astro-ph0112440_arXiv.txt": { + "abstract": "We have reconstructed the temperatures and surface densities in the quiescent accretion disc in HT~Cas by performing a {\\em Physical Parameter Eclipse Mapping} analysis of archival UBVR observations. Using a simple hydrogen slab model and demanding a smooth, maximally artefact-free reconstruction, we derive a formal distance to HT~Cas of $207\\pm 10$\\,pc, significantly larger than the $133\\pm 14$\\,pc we derive from a re-analysis of the data in the literature. The accretion disc is small ($0.3-0.4~\\Rl$) and moderately optically thin but becomes nearly optically thick near the white dwarf. The temperatures and surface densities in the disc range from 9\\,500\\,K and 0.013\\,g\\,cm$^{-2}$ in the center to about 4\\,000\\,K and 0.04\\,g\\,cm$^{-2}$ at the disc edge. The mass-accretion rate in the disc is roughly constant but -- at the derived distance -- uncomfortably close to those which would prohibit the dwarf nova eruptions. We argue that the larger derived distance is probably incorrect but is not produced by inaccuracies in our spectral model or optimization method. The discrepancy can be resolved if the emission regions on the disc are patchy with a filling factor of about 40\\% of the disc's surface. This solves the problem with the high effective temperatures in the disc -- reducing them to around 6\\,500\\,K within a radius of $0.2\\,\\Rl$ -- and reduces the derived temperature of the white dwarf and/or boundary layer from 22\\,600 to 15\\,500\\,K. The viscosity parameters $\\alpha$ derived from all reconstructed temperatures and surface densities are of order 10-100 and cannot be lowered significantly by invoking a lower distance or the filling factor. This situation is easily explained using the same patchy nature of the emitting material, since the quiescent disc cannot consist of optically thin regions alone, but also of a dark and hence cold and dense disc which could easily contain most of the matter. If we require global values of $\\alpha$ of order 0.1, the implied total surface densities are 1-100\\,\\gcm\\ -- just like those expected for quiescent discs awaiting the next eruption. We discuss several possible sources of the chromospheric emission and its patchiness, including irradiation of the disc, thermal instabilities, spiral-wave-like global structures, and magnetically active regions associated with dynamo action and/or Balbus-Hawley instabilities. ", + "introduction": "Cataclysmic variables (CVs), are semi-detached binary stars consisting of a Roche lobe-filling late-type dwarf (the secondary star) which loses matter onto a white dwarf (the primary). Unless the primary has a considerable magnetic field, the transfered matter has enough specific angular momentum to create a gaseous accretion disc around this accreting star. Some CVs -- the dwarf novae -- have discs which occasionally have luminous phases of high mass-accretion but spend most of their time in a faint quiescent state. In quiescence, the optical light is usually dominated by optically thin line emission from the disc and continuum emission from the white dwarf, the disc and the bright spot caused by the impact of the transferred material onto the disc. While spectral and eclipse observations of erupting dwarf novae show that the disc is optically thick during the outburst state, there is considerable uncertainty about the state of the quiescent discs. Theoretical models for the outbursts invoking a thermal instability as the cause of the eruption require that the quiescent disc store up material for the next eruption in a less viscous state in which the disc should be cool and optically thick (e.g. Ludwig et al. 1994). The inner disc may be emptied by irradiation from the hot white dwarf (Leach et al. 1999) or by a siphon flow fueled by a hot corona in the inner disc (Meyer \\& Meyer-Hofmeister 1994). Thus, the expectation is that the outer disc is cool and optically thick and the inner disc is warm/hot and optically thin. HT~Cas was once called ``the {\\em Rosetta Stone} of dwarf novae'', because -- being one of the relatively few eclipsing dwarf novae -- it more readily revealed its secrets (Patterson 1981). It has seldom been caught in outburst because of the unusually long outburst period of about 400 $\\pm$ 50 days (Wenzel 1987). The quiescent periods may even last up to almost 9 years\\footnote{Since it is a circumpolar object for many northern observatories, the likelihood of detecting outbursts is relatively high.} indicating an extremely low mass accretion rate in quiescence. The accretion disc is hardly visible during quiescence and the bright spot, the impact region of the accretion stream from the secondary, is nearly absent (Wood, Horne \\& Vennes et al.\\ 1992, hereafter WHV92). The light curves predominantly show the eclipse of the white dwarf. The optical spectrum of HT~Cas (Young, Schneider \\& Shectman 1981) shows double peaked emission lines of H~I, He~I, He~II and Ca~II, typical for a high inclination system ($i = 81^\\circ$: Horne, Wood \\& Stiening 1991, hereafter HWS91). Consistent with the broad-band photometry, the averaged profiles do not show any significant blue/red asymmetry, which would normally be attributed to the bright spot. WHV92 investigated the accretion disc using multicolour eclipse photometry from which the contribution of the white dwarf had been subtracted. They produced independent classical eclipse maps for each colour and then modeled them with a two-parameter model spectrum yielding the temperature and surface density distribution in the disc. While this analysis suggested that the disc is optically thin -- as expected -- it suffers in detail due to the individual smearing of each eclipse map (see Vrielmann, Horne \\& Hessman 1999; hereafter VHH99). Furthermore, the pre-subtraction of the white dwarf contribution and the use of simple Cartesian reconstruction maps can make it difficult to produce a consistent physical picture of the various components. \\begin{figure} \\psfig{file=ht205_175f.ps,width=8cm} \\caption{\\small The averaged quiescent light curves in UBVR together with the fits made simultaneously in all four filters. The contact phases of the white dwarf as given by Wood, Horne \\& Vennes (1992) are drawn as vertical dashed lines. The light curves are offset by 1.5 mJy, the horizontal dotted lines give the zero level for each light curve. In relation to this, the uneclipsed components and the residuals are plotted (solid lines). \\label{htf}} \\end{figure} Our present investigation is aimed at improving our understanding of the physical state of the quiescent accretion disc in HT~Cas. We apply the {\\it Physical Parameter Eclipse Mapping} method (PPEM) -- a distinctly different tomographic ansatz (VHH99) which also uses the Maximum-Entropy Method (MEM; Skilling \\& Bryan 1984) -- to the same UBVR photometry and derive the white dwarf temperature and the disc structure by simultaneously fitting all available eclipse light curves. ", + "conclusions": "We have applied the {Physical Parameter Eclipse Mapping} (PPEM) method to archival UBVR photometry of the eclipses in the quiescent dwarf nova HT~Cas in order to map the optically thin matter responsible for the optical continuum and emission line spectrum. Using a spectroscopically determined distance, we were unable to obtain reasonable maps of the temperature and surface density in the disc: the reconstructions showed characteristic artefacts showing that the surface brightness of the disc was too high. By insisting that the disc be maximally structureless, we were able to find reasonable solutions at a distance of 207pc. The resulting accretion disc has a radius of 0.3-0.4$\\Rl = 9-13 \\times 10^9$~cm, kinetic temperatures between $4000$~K and $10^4$~K, surface densities between 0.013 and 0.04\\,g\\,cm$^{-2}$, and $\\tau \\lsim 1$ everywhere except near the hot central white dwarf (22\\,600~K). The effective temperature is flat in the inner part of the disc and may follow a steady state solution towards the disc edge. We derive apparent values of the viscosity parameter $\\alpha$ from the bolometric fluxes, temperatures, and surface densities which are orders of magnitude higher than they should be. Acknowledging that the distance to HT~Cas is unlikely to be as high as suggested by our PPEM entropy analysis, we explain why the problem is unlikely to be due to a poor treatment of the white dwarf or the wrong choice of a spectral model. We are able to rectify both the poor reconstructions of the models using $d \\sim 150$~pc and the problems with those using $d \\sim 200$~pc by invoking emission regions which are not distributed uniformly across the face of the disc but are patches with a covering factor of only 41\\%. Since the angular extents of the emitting regions in the distant non-patchy and close by patchy discs can be identical, the reconstructed physical parameters of the $d \\sim 200$~pc solutions are still valid at $d \\sim 150$~pc with the exception of those quantities which depend upon the bolometric fluxes: the mass-accretion rates (or effective temperatures) need to be scaled down by 80\\%. This reduces the local mass-accretion rates to just at or below the critical rates for maintaining the dwarf nova eruptions for all but the innermost disc. The empirical values of $\\alpha$ (cf.~\\ref{viscosity}) -- ranging from about 10 to 1\\,000 -- for both $d \\sim 150$~pc and $d \\sim 200$~pc are still orders of magnitude larger than is physically plausible. Fortunately, the patchy chromosphere on the disc surface requires a substantial mass of very cool material which powers the viscous energy finally radiated away in the chromosphere and which have surface densities of the order of tens or hundreds of g\\,cm$^{-2}$ -- just those needed by the disc instability models. The patchiness excludes models for disc coronae due to local thermal instabilities or to irradiation. If we interpret the patchy chromospheric emission regions as magnetically active regions, the fact that the dissipation is proportional to the angular velocities could be a sign that what we are seeing magnetic flux created by magnetohydrodynamical instabilities and/or disc dynamos responsible for the anomalous viscosity in discs dissipated in magnetically active regions. However, the presence of ``mirror eclipses'' in another dwarf nova implies that the warm chromospheric patches may be situated laterally adjacent to rather than simply above the cold, more massive regions. HT~Cas may really deserve its name as ``Rosetta Stone'' of dwarf novae, since it could be the key to understanding the anomalous viscosity in accretion discs as due to hydromagnetic turbulence. High time and spectral resolution spectroscopy using a 10-m class telescope and/or multi-colour observations in a yet to be detected low state could provide the final proof for our hypothesis. The Physical Parameter Eclipse Mapping proves to be a powerful tool to investigate the physics of accretion discs -- even when it at first appears to produce a wrong result!" + }, + "0112/astro-ph0112489_arXiv.txt": { + "abstract": "The local expansion field is mapped using Cepheids, a complete sample of TF distances, and nearby cluster distances. The large-scale field is mapped using Cepheid-calibrated blue SNe\\,Ia. These data give $H_0({\\rm local})=59.2\\pm1.4\\;[\\mbox{km\\,s}^{-1}\\,\\mbox{Mpc}^{-1}]$ and $H_0({\\rm cosmic})=57.4\\pm2.3$. The intermediate expansion field ($1200\\le v \\la 10\\,000\\kms$) is less well calibrated but fully consistent with $H_0\\approx60$. $H_0$ is therefore (nearly) scale-invariant (high-density regions excluded). -- The P-L relation of Cepheids is based on an improved zero point of $(m-M)_{\\rm LMC}=18.56$. The slope of the P-L relation for $P>10^{\\rm d}$, as judged from OGLE data (Udalski et~al. 1999) is flatter than anticipated, which tends to increase the above values of $H_0$ by 3.4 units. No significant metallicity effect on the Cepheid distances seems to be indicated. For all practical purposes $H_0=60$ is recommended with a systematic error of probably less than 10\\%. -- The corresponding expansion age is $T=15.7\\pm1.5\\;$Gy (with $\\Omega_{\\rm m}=0.3$, $\\Omega_{\\Lambda}=0.7$), which compares well with the formation time of $15\\pm2\\;$Gy for the apparently oldest globular cluster M\\,107. ", + "introduction": "Cepheids and their $P$-$L$ relation are the principal fundament of extragalactic distances (Section~2). Cepheid distances are available of nine galaxies which have produced blue SNe\\,Ia (Saha et~al\\ 2001 and references therein). Their resulting, very uniform luminosities can be applied to 35 more distant SNe\\,Ia ($v \\la 30\\,000\\kms$) to yield a first rate determination of the large-scale value of $H_0$ (Section~5). Details of the local expansion field are inserted in Section~3 and some remarks on the not so local expansion field in Section~4. The expansion age of the Universe is compared with independent age determination of old objects in the Galaxy in Section~6. Results and conclusions are compiled in Section~7. ", + "conclusions": "\\label{sec:7} A synopsis of the distance determinations in Sections~3 and 5 is shown in Fig.~\\ref{fig:9}. The distance moduli of the SNe\\,Ia are determined here from their apparent magnitudes $m^{\\rm corr}_{B,V,I}$ and the absolute magnitudes $M^{\\rm corr}_{B,V,I}$ in equation~(\\ref{eq:13}). The local expansion field extends into the cosmic expansion without any significant break. If one choses to base the Cepheid distances of the SNe\\,Ia-calibrating galaxies on the new $P$-$L$ relations (eq.~\\ref{eq:7} \\& \\ref{eq:8}), they are reduced by 6\\% (cf.\\ eq.~\\ref{eq:8}), i.e. \\begin{equation}\\label{eq:16} H_0({\\rm cosmic})=60.8\\pm2.3. \\end{equation} For all practical applications $H_0=60$ can be used everywhere, except in nearby high-density regions. So far only statistical errors have been quoted. It comes as a surprise that the largest source of systematic errors is in the {\\em shape\\/} of the $P$-$L$ relation (6\\%), followed by the metallicity dependence of Cepheids and the photometric HST zero point in the crowded fields of SNe\\,Ia-calibrating galaxies (4\\%). The zero point of the $P$-$L$ relation, the slope of the $\\Delta m_{15}$ correction of SNe\\,Ia and the HST photometry may each contribute systematic 2-3\\% errors. Systematic errors due to absorption corrections for the nearby, calibrating SNe\\,Ia and the distant SNe\\,Ia are negligible, because the two sets have closely the same colors ($<\\!B-V\\!>=-0.01\\pm0.01$; cf.\\ Parodi et~al.\\ 2000). Unless there is a conspiracy of the individual systematic errors, the total systematic error is $<10\\%$. The resulting expansion age of $T=15.7\\pm1.5\\;$Gy ($H_0=60\\pm5$, $\\Omega_{\\rm m}=0.3$, $\\Omega_{\\Lambda}=0.7$) gives sufficient room for the oldest dated objects in the Galaxy. \\smallskip\\noindent {\\bf Acknowledgments.} The first three authors thank the Swiss National Science Foundation for financial support. G.\\,A.\\,T. and F.\\,T. thank also the PRODEX programme of the Swiss Space Office for support. \\footnotesize" + }, + "0112/astro-ph0112395_arXiv.txt": { + "abstract": "In view of the extensive evidence of tight inter-relationships between spheroidal galaxies (and galactic bulges) with massive black holes hosted at their centers, a consistent model must deal jointly with the evolution of the two components. We describe one such model, which successfully accounts for the local luminosity function of spheroidal galaxies, for their photometric and chemical properties, for deep galaxy counts in different wavebands, including those in the (sub)-mm region which proved to be critical for current semi-analytic models stemming from the standard hierarchical clustering picture, for clustering properties of SCUBA galaxies, of EROs, and of LBGs, as well as for the local mass function of massive black holes and for quasar evolution. Predictions that can be tested by surveys carried out by SIRTF are presented. ", + "introduction": "The hierarchical clustering model with a scale invariant spectrum of density perturbations in a Cold Dark Matter (CDM) dominated universe has proven to be remarkably successful in matching the observed large-scale structure as well as a broad variety of properties of galaxies of the different morphological types (Granato et al. 2000 and references therein). Serious shortcomings of this scenario have also become evident in recent years. At the other extreme of the galaxy mass function with respect to so-called ``small-scale crisis'', another strong discrepancy with model predictions arises, that we might call ``the massive galaxy crisis''. Even the best semi-analytic models hinging upon the standard picture for structure formation in the framework of the hierarchical clustering paradigm, are stubbornly unable to account for the (sub)-mm (SCUBA, see Fig.~1, and MAMBO) counts of galaxies, most of which are probably massive objects undergoing a very intense star-burst (with star formation rates $\\sim 1000\\,\\hbox{M}_\\odot\\,\\hbox{yr}^{-1}$) at $z>2$. Recent optical data confirm that most massive ellipticals were already in place and (almost) passively evolving up to $z\\simeq 1$--1.5. These data are more consistent with the traditional ``monolithic'' approach whereby giant ellipticals formed most of their stars in a single gigantic starburst at substantial redshifts, an underwent essentially passive evolution thereafter. In the canonical hierarchical clustering paradigm the smallest objects collapse first and most star formation occurs, at relatively low rates, within relatively small proto-galaxies, that later merged to form larger galaxies. Thus the expected number of galaxies with very intense star formation is far less than detected in SCUBA and MAMBO surveys and the surface density of massive evolved ellipticals at $z\\gsim 1$ is also smaller that observed. The ``monolithic'' approach, however, is inadequate to the extent that it cannot be fitted in a consistent scenario for structure formation from primordial density fluctuations. ", + "conclusions": "" + }, + "0112/astro-ph0112506_arXiv.txt": { + "abstract": "We analyze observations of the microwave sky made with the Python experiment in its fifth year of operation at the Amundsen-Scott South Pole Station in Antarctica. After modeling the noise and constructing a map, we extract the cosmic signal from the data. We simultaneously estimate the angular power spectrum in eight bands ranging from large ($\\ell \\sim 40$) to small ($\\ell \\sim 260$) angular scales, with power detected in the first six bands. There is a significant rise in the power spectrum from large to smaller ($\\ell \\sim 200$) scales, consistent with that expected from acoustic oscillations in the early Universe. We compare this Python V map to a map made from data taken in the third year of Python. Python III observations were made at a frequency of 90 GHz and covered a subset of the region of the sky covered by Python V observations, which were made at 40 GHz. Good agreement is obtained both visually (with a filtered version of the map) and via a likelihood ratio test. ", + "introduction": "Since the detection of anisotropy in the cosmic microwave background by the COBE satellite, many experiments have measured the angular power spectrum at degree and sub-degree angular scales (e.g., Netterfield et al. 2002; Halverson et al. 2002; Lee et al. 2001; Miller et al. 1999). The Python V data set has sufficient sky coverage to probe the smallest scales to which COBE was sensitive, while having a small enough beam to detect the rise in the angular power spectrum to degree angular scales, providing a link in $\\ell$-space between COBE and other recent measurements. Python V is the latest of the Python experiments at the South Pole. Dragovan et al. (1994), Ruhl et al. (1995), and Platt et al. (1997) describe Python I--III and Rocha, et al. (1999) derive constraints on cosmological parameters from these data. Kovac et al. (1997) describe the Python IV results. The Python V experiment, observations, and data reduction are described in Coble et al. (1999). In that paper, we analyzed individual modulations of the data. The modulations can be thought of as filters which have little sensitivity to some of the contaminants in the time stream. For example, they have no sensitivity to gradients, which should get a large contribution from the atmosphere and from the ground shield. The modulation approach also provided a rapid means of compressing a large amount of data (19 Gbytes) into a more manageable size. Measurements of anisotropy were reported for eight different modulations of the sky signal; the results indicated a sharp rise in the power spectrum. In this paper we find the constraints on the power spectrum due to all of the modulations simultaneously. We use the modulations as our starting point, rather than the time stream, to take advantage of the contaminant filtering and data compression. We extend the analysis of Coble et al. (1999) by accounting for the correlations (in both signal and noise) between different modulations. From the modulations we find the best-fit map and its associated noise covariance. From this map and its associated covariance matrix we estimate the power spectrum simultaneously in eight bands. In $\\S$ 2 we briefly review the instrument and the data set. In $\\S$ 3 we discuss the estimation of the noise matrix. In $\\S$ 4 we describe how to use this matrix to construct a map and a noise matrix for the map. This map is used in $\\S$ 5 to estimate the angular power spectrum in eight bands. In $\\S$ 6 we check the power spectrum derived from the map with the power spectrum derived directly from the modulated data. In $\\S$ 7 we compare the 40 GHz Python V data with the 90 GHz Python III data (Platt et al. 1997), which covered a subset of the region of the sky covered by Python V. We find good agreement between the two observations in the region of overlap, providing a valuable consistency check. This is another indication of a lack of significant foreground contamination (see also our estimates in Ganga et al. 2002). We conclude in $\\S$ 8. ", + "conclusions": "The Python V experiment densely samples 598 square degrees of the microwave sky and constrains the CMB anisotropy angular power spectrum from $\\ell \\sim$ 40 to $\\ell \\sim$ 260, showing that power is increasing from large to smaller ($\\ell \\sim$ 200) angular scales. The noise matrix constructed in $\\S$ 3 enables us to simultaneously estimate the angular power spectrum in eight bands. The power spectra estimated from the map and directly from the modulated data are consistent. The rise seen in Figure \\ref{fig6} is characteristic of acoustic oscillations in the early Universe. A number of other measurements also indicate such a rise in power (e.g., Ganga, Ratra, \\& Sugiyama 1996; Netterfield et al. 1997; de Oliveira-Costa et al. 1998; Torbet et al. 1999; Podariu et al. 2001; as well as experiments mentioned in $\\S$~1). Python V extends to larger scales (lower $\\ell$) than these, to the smallest scales to which COBE was sensitive. The Python III and V experiments differ in significant ways, including frequency, receiver, year, and noise properties. Nevertheless, the maps and the $\\beta$ test in $\\S$ 7 indicate that they both detect similar signals, a rare and very valuable consistency check and confirmation." + }, + "0112/nucl-th0112032_arXiv.txt": { + "abstract": " ", + "introduction": "Collapse driven supernovae (SNe) and the following formation of neutron stars (NSs) are the most dramatic processes during the stellar evolution. These objects provide not only astrophysically significant phenomena but also interesting material phases inside them; both are strongly connected with each other. At subnuclear densities in NS crusts and SN cores, nuclei are expected to have exotic structures such as rod-like and slab-like shape etc., which are called nuclear ``pasta''.\\cite{rf:review} This prediction has been obtained by previous studies \\cite{rf:rpw,rf:hashimoto} based on free energy calculations assuming specific nuclear shapes. These works clarify that the nuclear shape is determined by a subtle balance between nuclear surface and Coulomb energies. However, the formation process of the ``pasta'' phases has not been discussed except for some limited cases which are based on perturbative approaches. \\cite{rf:review,rf:iida} It is important to adopt microscopic and dynamical approach which allows arbitrary nuclear structures in order to understand the formation of the non-spherical nuclei. We have started studying dense matter by the QMD\\cite{rf:qmd} simulation whose final goal is to obtain a realistic picture about NS crusts and SN cores. As the first step of our quest, we will answer the following question in this study: ``Can the 'pasta' phases be formed dynamically?'' ", + "conclusions": "" + }, + "0112/astro-ph0112339_arXiv.txt": { + "abstract": "We present ISOPHOT spectrophotometric observations of SS\\,433 at four different orbital phases in 1996 and 1997. The He{\\footnotesize I}\\,+\\,He{\\footnotesize II} lines\\ \\ in\\ \\ both\\ \\ spectra\\ \\ of\\ \\ SS\\,433\\ \\ and\\ \\ the\\ \\ Wolf-Rayet\\ \\ star\\ \\ WR\\,147,\\ \\ a \\mbox{WN8+BO5} binary system, closely match. The 2.5\\,--\\,12\\,$\\mu$m continuum radiation is due to an expanding wind free-free emission in an intermediate case between optically thick and optically thin regime. A rough mass loss rate evaluation gives about $\\sim 1.4 \\times 10^{-4}\\,\\mathrm{M}_\\odot.\\mathrm{yr}^{-1}$. Results are consistent with a Wolf-Rayet-like companion to the compact object in SS\\,433. ", + "introduction": "The X-ray binary source SS\\,433 and its relativistic jets have been studied at many wavelengths, with yet no consensus about the component masses and the evolutionary status of the system (see Margon 1984 for a review). As shown by van den Heuvel, Ostriker, \\& Petterson\\,(1980) the optical stationary spectrum and photometric characteristics of SS\\,433 are consistent with those of an Of or Wolf-Rayet (WR)-like star with an extremely large rate of stellar wind mass loss. The predicted mass transfer rates are much larger than the likely mass loss rates in the precessing jets (Begelman et al. 1980). King, Taam, \\& Begelman\\,(2000) argue that most of the transferred mass is lost from the accretion flow at large radii and is presumably the source of the stationary H$\\alpha$ line and the associated free-free continuum in the near-infrared seen by Giles et al.\\,(1979). We\\ \\ present\\ \\ middle\\ \\ and\\ \\ far-infrared\\ \\ spectrophotometric\\ \\ observations\\ \\ of\\ \\ SS\\,433 with ISOPHOT (Kessler et al. 1996; Lemke et al. 1996) to test this interpretation. We will discuss the spectral shape of the continuum and compare the emission lines in SS\\,433 and the binary Wolf-Rayet star WR\\,147. ", + "conclusions": "We have shown that the mid-infrared continuum of SS\\,433 between 2.5 - 12 $\\mu$m can be explained by the free-free emission of an expanding wind in the intermediate case between optically thick and optically thin regime. This result is consistent with the assumed large mass expelled from the super-Eddington accretion disk at large radii (King et al. 2000), and with the wind-like equatorial outflow observed by Paragi et al.\\,(1999) - and their origin is probably a stellar wind from the companion star. The close match between the HeI + HeII emission lines detected in SS\\,433 and WR\\,147 is consistent with a Wolf-Rayet-like companion to the compact object. Thus SS\\,433 might be the second known X-ray binary containing a Wolf-Rayet star with a compact object after the galactic relativistic jet source Cygnus X-3 (Hanson, Still, \\& Fender 2000 and references therein). \\vspace*{-0.3cm} \\begin{figure}[!htp] \\centerline{\\psfig{figure=fuchsy_3.ps,angle=-90,width=10cm}} \\vspace*{-0.4cm} \\caption{Radio to optical spectrum of SS\\,433} \\label{fig2} \\end{figure}" + }, + "0112/astro-ph0112049_arXiv.txt": { + "abstract": "We examine Cosmic Microwave Background (CMB) temperature power spectra from the BOOMERANG, MAXIMA, and DASI experiments. We non-parametrically estimate the true power spectrum with no model assumptions. This is a significant departure from previous research which used either cosmological models or some other parameterized form (e.g. parabolic fits). Our non-parametric estimate is practically indistinguishable from the best fit cosmological model, thus lending independent support to the underlying physics that governs these models. We also generate a confidence set for the non-parametric fit and extract confidence intervals for the numbers, locations, and heights of peaks and the successive peak-to-peak height ratios. At the 95\\%, 68\\%, and 40\\% confidence levels, we find functions that fit the data with one, two, and three peaks respectively ($ 0 \\le \\ell \\le 1100$). Therefore, the current data prefer two peaks at the $1\\sigma$ level. However, we also rule out a constant temperature function at the $> 8 \\sigma$ level. If we assume that there are three peaks in the data, we find their locations to be within $\\ell_1$ = (118,300), $\\ell_2$ = (377,650), and $\\ell_3$ = (597,900). We find the ratio of the first peak-height to the second $(\\frac{\\Delta T_1}{\\Delta T_2})^2 = (1.06, 4.27) $ and the second to the third $(\\frac{\\Delta T_2}{\\Delta T_3})^2 = (0.41, 2.5)$. All measurements are for 95\\% confidence. If the standard errors on the temperature measurements were reduced to a third of what they are currently, as we expect to be achieved by the MAP and Planck CMB experiments, we could eliminate two-peak models at the 95\\% confidence limit. The non-parametric methodology discussed in this paper has many astrophysical applications. ", + "introduction": "There has been growing evidence for the existence of peaks and valleys in the temperature power spectrum of the CMB. \\relax From a theoretical standpoint, such features are a direct result of the physics in the primordial photon-electron plasma, predicted by gravitational instability models of structure formation (Peebles \\& Yu 1970; Hu and Sugiyama 1995). These features are important for constraining the cosmology of our Universe. For instance, in many models, the ratio of the height of the first peak to the second peak is dependent on the spectral tilt, $n_s$ and the baryon fraction, $\\Omega_{baryons}/\\Omega_{matter}$. The ratio of the third peak to the second peak is dependent on $\\Omega_{matter}h^2$ and $n_s$ (see Hu et al. 2001 for further discussion). Most often in the literature, the CMB power spectra are fit to a suite of cosmological models (Tegmark et al., 1999,2000; Jaffe et al. 2001). These physical models are well-motivated and sophisticated, but they contain many free parameters (\\emph{e.g.}, eleven in the work of Wang, Tegmark, and Zaldarriaga 2001-- WTZ), some of which are unknown (ionization depth, contribution from gravity waves) or degenerate (e.g. see Efstathiou 2001). Typically, some sort of likelihood analysis is performed to determine which cosmological model best fits the data. There is however, another approach: place constraints on the features of the power spectrum and use these features to determine the cosmological parameters. The assumptions here are that the peaks and valleys are best described by the broad range of cosmological models (as in Hu et al.) or by parabolas or some other chosen function (as in Knox \\& Page 2000 and de Bernardis et al. 2001). A potential problem in all of these approaches is that it is difficult to get valid statistical confidence intervals (see Section 2). There is also the concern that the fitted features may be artifacts from the multitude of assumptions. In this paper, we take what may be considered a more conservative approach: we make no assumptions whatsoever about the true underlying function. Our new statistical technique is non-parametric and allows for valid confidence intervals to be measured for peak characteristics. One theme of our work is that confidence intervals for any quantity of interest can be extracted from a confidence set for the unknown spectrum. These techniques for fitting and inference are applicable to a wide variety of astrophysical data-analysis problems. ", + "conclusions": "We present an application of a new and powerful non-parametric technique to CMB temperature data. Past approaches were based on complicated cosmological models or parameterized forms. There is superb visual agreement between the non-parametric fit and the best fitting cosmological model. Quantitatively, we provide constraints on the peak locations, heights, and height ratios of the power spectrum. These constraints can be used to place corresponding limits on the cosmological parameters that they describe. For instance, Hu et al. (2001) derive relationships which could in principle, be used for this purpose. At the $2\\sigma$ confidence level, we find at least one peak in the current CMB power spectra data, while at the $1 \\sigma$ level, we find two or more peaks. Only for a very low confidence, 40\\%, can we rule out two peak functions. Therefore, the data do not yet show the three expected peaks for $\\ell \\le 1100$ (in the three CMB datasets examined here). There are two explanations for this: the model is right, but there is insufficient precision in the current data, or the model is wrong. If in fact the errors on the current measurements are simply too large, then these standard errors would have to be reduced to one-third of their current values to rule out a two-peak spectrum at the 95\\% confidence level. This suggests a range for the maximum required errors for future CMB experiments (via MAP and Planck) to ``discover'' three peaks in the CMB spectrum. We point out that the lack of assumptions used to arrive at our best fit is conservative. On the other hand, results from fitting assumed cosmological models are optimistic, since those models all have a multi-peaked spectra (e.g. Hu et al. 2001). While the physical underpinnings for cosmological models are well founded, the last 50 years (or even five) have seen radical changes in those models which best fit the data. Therefore, a method to describe the CMB that is ``cosmology free'' has scientific value. Finally, we note that the methods described here can be applied to the many astrophysical problems that are not well suited for standard parametric techniques. \\vspace{0.2in} This work was done in collaboration with the Pittsburgh Computational Astrostatistics Group (\\verb+www.picagroup.org+). The authors would like to thank the referee for suggestions which improved the readability and usefulness of this work. RCN, LW, and CG were partially supported by NSF KDI grant DMS-9873442. \\vspace{0.2in} During the refereeing process of our paper, two related papers came to our attention (Durrer, Novosyadlyj \\& Apunevych 2001; Douspis \\& Ferreira 2001). These papers perform a model--independent measurement of the CMB power spectrum but they are not non-parametric estimates of the CMB acoustic peaks, as discussed herein, since they use phenomenological models to describe the underlying power spectrum. It is interesting to note however, that all three analyses find low statistical significances for the detection of the second and third peaks. We await higher precision measurements of the CMB power spectrum to secure the detection, location, shape and amplitude of these peaks." + }, + "0112/astro-ph0112563_arXiv.txt": { + "abstract": "In cosmological models with a varying gravitational constant, it is not clear whether primordial black holes preserve the value of $G$ at their formation epoch. We investigate this question by using the Tolman-Bondi model to study the evolution of a background scalar field when a black hole forms from the collapse of dust in a flat Friedmann universe. Providing the back reaction of the scalar field on the metric can be neglected, we find that the value of the scalar field at the event horizon very quickly assumes the background cosmological value. This suggests that there is very little gravitational memory. ", + "introduction": "Scalar-tensor (ST) theories of gravity provide a natural alternative to general relativity (GR). They describe gravity with not only a metric $g_{ab}$ but also a scalar field $\\phi$. Derivatives of $\\phi$ appear as source terms in the field equations and $\\phi$ itself satisfies a wave equation. The strength of the gravitational coupling is determined by the function $\\omega(\\phi)$, where GR is recovered in the limits $\\omega\\rightarrow\\infty$ and $\\omega^{-3}(d\\omega/d\\phi)\\rightarrow0$. ST theories can also be regarded as being equivalent to GR with a varying gravitational ``constant'' $G$. The most simple example of such a ST theory is Brans-Dicke theory \\cite{BD}, where $\\omega(\\phi)$ is constant and $G\\propto\\phi^{-1}$. However, weak field experiments have shown that $\\omega>500$ \\cite{Will} and so the deviation from GR is small. For more general ST theories, where $\\omega$ is not constant, it is possible that $\\omega$ was much smaller at earlier times. So observations allow such theories to greatly deviate from GR in the early universe. There has been a renewed interest in ST theories in recent years due to the effective low energy actions of string theory involving one or more scalar fields. These scalar fields enter the field equations in much the same way as the scalar field in ST theories \\cite{Call}. Also, the increasing popularity of inflation and quintessence suggests that scalar fields might need to be incorporated into cosmological models. The purpose of this paper is to study the effect of an evolving scalar field on the formation and evolution of a primordial black hole. In an asymptotically flat spacetime it is well known that a black hole radiates away any inhomogeneities in the scalar field until it becomes a stationary solution with constant $\\phi$ \\cite{Hawk}. This is a consequence of the famous ``no hair'' theorem. However, in ST cosmological models the scalar field is evolving with time and this would modify how the black hole evolves during its lifetime. Barrow \\cite{Barr} was the first to examine this problem. He considered the two extreme possibilities: scenario (A), where the scalar field evolves everywhere homogeneously in the same way as the cosmological background; and scenario (B), where the black hole forces the scalar field to remain constant in some local region around it. The second scenario Barrow called gravitational memory because the black hole would locally preserve the value of the scalar field from when it formed. Barrow \\& Carr \\cite{BC} studied the evolution of primordial black holes for these scenarios and found that either case results in a significant deviation from the usual GR analysis. They considered ST theories, where $G(\\phi)\\propto\\phi^{-1}$ as in the Brans-Dicke case. Since $\\phi$ increases with cosmological time, this implies that $G(\\phi)$ decreases, so black holes would take longer to radiate away their mass via Hawking evaporation than in the GR case. In both scenarios the black holes form when gravity is stronger, so the rate of evaporation is less, but in scenario (B) the strength never decreases and so the lifetime is even longer. The above scenarios are the two extremes and the reality is probably somewhere in between. Two more general scenarios have been proposed \\cite{GC}. In scenario (C) the scalar field evolves faster at the event horizon (EH) than at the particle horizon (PH), so $\\dot{\\phi}_{EH}>\\dot{\\phi}_{PH}$. Eventually the black hole must reach a stage where $\\dot{\\phi}_{EH}=\\dot{\\phi}_{PH}$, but this does not necessarily mean the scalar field is homogeneous since there could be some lag between the asymptotic and local increase. In scenario (D), the scalar field is evolving locally but at a slower rate than asymptotically, so $\\dot{\\phi}_{EH}<\\dot{\\phi}_{PH}$. There is still some gravitational memory but not in the strict sense of scenario (B). In this scenario the gradient of the scalar field is increasing but one would expect there to be some limit, due to the influx of scalar gravitational waves. Gravitational memory for black holes which are small compared to the cosmological scale (i.e. the particle horizon) has already been investigated by Jacobson \\cite{Jaco}. In this case, the scalar field evolution can be considered as an asymptotic perturbation to the Schwarzschild metric and the end-state of scenario (C) applies, with the lag being found to be small. This suggests that gravitational memory is virtually non-existent. However, this approximation may not apply for primordial black holes since these can have a size comparable to the particle horizon at formation \\cite{CH}. It is still not clear what would happen in this case. Therefore another way of investigating gravitational memory, without assuming that the black hole is small, is needed. In GR there have been several attempts to study analytic solutions which represent black holes within a cosmological background. The earliest used the Einstein-Straus solution \\cite{ES}, which matches a Schwarzschild interior to a Friedmann exterior, and this approach has also been used to study gravitational memory \\cite{Sak}. However, in most circumstances it can be shown \\cite{Whin} that such a matching is only possible if the scalar field is constant, which just gives the GR solution. Another method used the McVittie metric \\cite{McV}, but it has been shown \\cite{Nola} that this has a scalar curvature singularity at the event horizon. A more successful method uses the Tolman-Bondi metric \\cite{TB} to represent the collapse of dust to a black hole in an asymptotically Friedmann background. However, this only works for dust and cannot in general be applied to ST theories due to the derivatives of the scalar field appearing as source terms in the field equations. In this paper we overcome this problem by assuming that the effect of the scalar field on the spacetime is small compared to that of the matter. This means that we can use the usual Einstein field equations to generate the spacetime and then use the wave equation for the scalar field to determine its evolution. This approximation was used by Harada et al. \\cite{HCNN} to calculate the scalar gravitational radiation emitted by Oppenheimer-Snyder collapse in ST theory. Jacobson also uses this approximation when calculating the effect of an evolving scalar field in Schwarzschild spacetime. We note, however, that self-consistent numerical calculations of spherical gravitational collapse in ST theory in asymptotic flat spacetime, in which the effect of the scalar field on the spacetime is fully incorporated, have been considered by previous authors \\cite{SNN,SST}. If one makes this approximation to investigate gravitational collapse in a Tolman-Bondi spacetime, the solution is specified by two arbitrary functions: the energy and mass functions. To represent collapse in a Friedmann background, the energy function has to be negative within some radius $r_0$ and zero outside it. This results in the eventual gravitational collapse of all the matter within $r_0$, while the matter outside $r_0$ expands forever as in a flat Friedmann model. In choosing the mass function, or equivalently the density perturbation, we adopt the ``compensated\" model. This means that the overdense region in the centre is surrounded by an underdense region outside, so that the total mass at infinity is unaffected. We solve the field equations numerically using the characteristic method. This method was first applied to an inhomogeneous and dynamical background spacetime by Iguchi et al. \\cite{IHN}. The characteristic method integrates over null hypersurfaces with the event horizon as a boundary. This means that one never needs to calculate anything inside the black hole, thereby avoiding any numerical problems associated with singularities. The output of the code shows the spatial and temporal variation of the scalar field. The figures produced show that the initial collapse results in a large gradient in the scalar field. However, as time increases, the scalar field becomes almost homogeneous. This suggests that, within the approximation used, gravitational memory is not possible. It remains to be seen whether the back reaction of the scalar field could alter this conclusion. In Section II we describe ST theories in more detail and derive the field equations for the approximation in which the effect of the scalar field on the metric is neglected. In Section III we transform the Tolman-Bondi solution into null coordinates, giving the equations necessary to apply the characteristic method. In Section IV we specify the model giving rise to black hole formation. We present the numerical results in Section V and discuss their implications in Section VI. ", + "conclusions": "We have calculated the evolution of the Brans-Dicke scalar field in the presence of a primordial black hole formed in a flat Friedmann background. We have found that the value of the scalar field at the event horizon almost always maintains the cosmological value. This suggests that primordial black holes ``forget'' the value of the gravitational constant at their formation epoch. In this sense, we confirm the result of Jacobsen \\cite{Jaco}, although he never carried out an explicit evolutionary calculation. While it is possible that some choices of the function $f(r)$ might lead to a different conclusion, we would claim that our conclusion must at least hold in compensated models. For we have tried many sets of parameters, corresponding to models which describe a wide variety of physical situations, and this always seems to be the case. However, it is not clear from our analysis what would happen for non-compensated models and we have also avoided models with shell-crossing singularities, so we do not claim that our result is completely general. It should also be stressed that this result has only been demonstrated for a dust universe in which the scalar field does not appreciably affect the background curvature and it remains to be seen whether the same conclusion applies when this assumption is dropped. As can be seen from Figs. \\ref{fg:tp} and \\ref{fg:rp}, both the radial gradient and the time derivative of $\\phi$ are large at early times and both of these act as source terms in the field equations. However, as long as $\\omega$ is large, this is unlikely to stop the scalar field becoming homogeneous eventually, since it is clear that both the radial gradient and time derivative become small well after the initial collapse. Neglecting the back reaction should therefore be a reasonable assumption in this situation, so gravitational memory seems unlikely. However, this conclusion might not apply for $\\omega \\approx 1$." + }, + "0112/astro-ph0112080_arXiv.txt": { + "abstract": "We still do not have a ``standard\" model for the mass distribution of spiral galaxies. I review various methods used to delimit the range of values of the mass-to-light ratio of the disk, such as spiral structure criteria, the influence of bars, warped and flaring HI layers, velocity dispersions, and features in the rotation curve occuring just beyond the optical radius. It is not yet clear whether disks are maximal or not. ", + "introduction": "In his classic paper on exponential disks (Freeman, 1970), Ken presents a unifying scheme to describe the photometric properties of disk galaxies, using a simple law which describes the radial distribution of the surface brightness of the disk of a spiral or S0 galaxy reasonably well. He then works out the dynamics of such disks, and discusses issues related to disk formation. In the Appendix of that paper he concentrates on four galaxies, and points out that for two of them the available data hint at a discrepancy between the photometric scalelength of the disk and the expected peak of the rotation curve if the mass-to-light ratio of the disk is independent of radius. While for the LMC and the SMC there is no such discrepancy, the available data for M33 and NGC 300 do indicate one: {``For two of four systems with weak spheroidal components, the LMC and the SMC, the turnover radii of the rotation curves predicted from the photometrically derived gradients $\\alpha$ appear to be consistent with those observed. For NGC 300 and M33, the 21-cm data give turnover points near the photometric edges of these systems. These data have relatively low spatial resolution; it they are correct, then there must be in these galaxies additional matter which is undetected, either optically or at 21 cm. Its mass must be at least as large as the mass of the detected galaxy, and its distribution must be quite different from the exponential distribution which holds for the optical galaxy.''} These remarks have been quoted by several authors retracing the history of the dark matter problem in spiral galaxies (e.g. Van Albada et al. 1985, Sciama, 1993). It is actually interesting to examine how the interpretation of the data of these two galaxies changed as the instrumentation improved. This is of intellectual relevance to see how scientific debate advances with time : it is sometimes said that the dark matter problem constitutes a ``revolution'', in the sense of Kuhn's book (Kuhn 1962), or worse a ``paradigm shift'' -- but then ``what is a paradigm ?'' (cf. Lakatos \\& Musgrave 1970) -- so what is the contribution of Ken's paper to this revolution ? Early data on M33 are very confusing, and the most up-to-date information available in 1969 is summarized in an IAU Symposium published in Russian (cf. Demoulin et al. 1969). Dramatic improvements in the kinematical data on M33 were obtained in the early 70's, first through the work of the Cambridge HI group (Warner et al. 1973), and then by Rogstad et al. (1976). In the latter paper, one of the key characteristics of the HI distribution was pointed out : the HI layer is strongly warped beyond the edge of the optical disk. It turns out that if the points on the rotation curve in the warped part are ignored, as done by Kalnajs (1983), one can fit the rotation curve assuming the mass follows the light. If the rotation curve is extended with the HI points from the warped parts, this cannot be achieved anymore (Athanassoula et al. 1987, fig.~5). For NGC 300, the early HI data are from a Parkes single dish study, (Shobbrook \\& Robinson 1967), and the surface photometry from De Vaucouleurs \\& Page (1962). Ken pointed out the discrepancy between the expected and observed turnover point of the rotation curve. HI data with much better angular resolution, analysed by Rogstad et al. (1979), show that the HI layer is strongly warped beyond the optical disk. Further study of this galaxy was done by a Puche, Carignan \\& Bosma (1990) -- a student of a student of Ken, a student of Ken and a postdoc of Ken -- which showed that there is indeed a mass discrepancy starting well within the optical radius, vindicating Ken's 1970 remark. Thus we can say that Ken's study, as he reported it in the Appendix of his 1970 paper, pointed for the first time in an elegant manner to a mass discrepancy in a galaxy, based on a then novel general description - which is still standing - of how the luminosity is distributed in the disk (the exponential law), the idea that the mass-to-light ratio in a disk could well be constant, and appealing to HI studies to get information on galaxy rotation curves. Dark halos came later (Ostriker \\& Peebles 1973), and their implication that spiral galaxy masses could really be large a year after (Einasto et al. 1974, Ostriker et al. 1974). These studies based themselves on the early work by Roberts \\& Rots (1973) and Rogstad \\& Shostak (1972). Yet at the observational side a long debate was spun about the validity of the data (in particular for M31, e.g. Roberts 1975, Baldwin 1975), the question whether fits with constant M/L were still possible (Emerson \\& Baldwin 1973, Baldwin 1975), the technical questions of sidelobes and beam smearing (e.g. discussions by Sancisi and Bosma following Salpeter 1977), and the interpretation of rotation curves derived from tilted ring models to describe warps (cf. discussion after Toomre 1977). Many of these problems were settled using high resolution HI data from the Westerbork telescope, accompanied by surface photometry and H$\\alpha$ kinematics (Bosma 1978, 1981, Bosma \\& Van der Kruit 1979, Van der Kruit \\& Bosma 1978). The work of Rubin et al. (1978), which later led to a series of papers on the rotation curves of Sc, Sb and Sa galaxies (Rubin et al. 1980, 1982, and 1985) was contemporary, but the HI data extends further out. Moreover, when, in the early 80s, everybody thought that dark halos are necessary to explain the flatness of rotation curves, Kalnajs (1983) showed that for optical rotation data, in four cases, there is no need for a dark halo, a conclusion corroborated for many more optical rotation curves by Kent (1986, 1987), Athanassoula et al. (1987), Buchhorn (1992) as reported in Freeman (1992), Moriondo et al. (1999), and Palunas \\& Williams (2000). ", + "conclusions": "A comparison with Bosma (1999) and Sellwood (1999) shows where progress has been made on the issue of whether disks are self-gravitating: the results from Athanassoula et al. (1987) based on the swing amplifier criteria still stand, new results from Kranz et al. (2001) favour less than maximum disks, several studies argue that barred galaxies have maximum disks, and the use of stellar velocity dispersions is hampered by lack of good data. So not much progress ? Better data will come from the gravitational lensing methods, which rather soon might come up with a clear answer. Until then, we can still try to improve things by getting better data and simulations along the lines discussed above." + }, + "0112/astro-ph0112033_arXiv.txt": { + "abstract": "We present results from the observation of the young Galactic supernova remnant (SNR) G292.0+1.8 with the Advanced CCD Imaging Spectrometer (ACIS) on board the {\\it Chandra X-ray Observatory}. In the 0.3 $-$ 8 keV band, the high resolution ACIS images reveal a complex morphology consisting of knots and filaments, as well as the blast wave around the periphery of the SNR. We present equivalent width (EW) maps for the elemental species O, Ne, Mg, and Si, which allow us to identify regions of enhanced metallicity in the SNR. G292.0+1.8 is bright in O, Ne, and Si; weaker in S and Ar; with little Fe. The EW and broad-band images indicate that the metal-rich ejecta are distributed primarily around the periphery of the SNR. The central belt-like structure has normal solar-type composition, strongly suggesting that it is primarily emission from shocked circumstellar medium rather than metal-rich ejecta. We propose that the belt traces its origin to enhanced mass loss in the star's equatorial plane during the slow, red supergiant phase. We also identify thin filaments with normal composition, centered on and extending nearly continuously around the outer boundary of the SNR. These may originate in a shell caused by the stellar winds from the massive progenitor in the red/blue supergiant phases, over-run by the blast wave. ", + "introduction": "OBSERVATION \\& DATA REDUCTION} G292.0+1.8 was observed with the Advanced CCD Imaging Spectrometer (ACIS) on board {\\it Chandra} on 2000 March 11 as part of the {\\it Chandra} Guaranteed Time Observation program. The angular size of G292.0+1.8 is slightly larger than the 8$'$ size of a single CCD chip; the ACIS-S3 chip was chosen because it has the best energy resolution over the entire CCD. The pointing was selected to place most of the SNR on the S3 chip. We have utilized new data reduction techniques developed at Penn State for correcting the spatial and spectral degradation of ACIS data caused by radiation damage, known as Charge Transfer Inefficiency (CTI) \\citep{townsley00}. The expected effects of the CTI correction include an increase in the number of detected events and improved event energies and energy resolution \\citep{townsley00,townsley01}. We screened the data with the flight timeline filter and then applied the CTI correction before further data screenings by status, grade, and the energy selections. ``Flaring'' pixels were removed and {\\it ASCA} grades (02346) were selected. Events with inferred energies between 0.3 keV and 8.0 keV were extracted for further data analysis. The overall lightcurve was examined for possible contamination from time variable background and no significant variability was found. After data processing $\\sim$43~ks of effective exposure was left. ", + "conclusions": "SUMMARY} We have presented results from the {\\it Chandra}/ACIS observation of the Galactic O-rich SNR, G292.0+1.8. The high resolution ACIS images resolve the complex filamentary structure in the soft band while revealing a pulsar and its wind nebula in the hard band. The EW maps % indicate variable elemental abundance and ionization across the SNR. Based on the surface brightness images and the EW maps, we propose that the equatorial belt in the broad band images is emission from the shocked non-spherically symmetric CSM, produced by slow stellar winds from the massive progenitor in the red supergiant phase, rather than from metal-rich ejecta. We also suggest that the origin of the soft shell-like structure around the outer boundary of the SNR may be emission from the shell of the stellar winds over-run by the blast wave, or the shock front itself if the X-ray emission of larger radii comes from portions of the remnant seen in projection along the line of sight. The distribution of the ejecta implies a complex evolution for G292.0+1.8, such as an asymmetric SN explosion and highly clumped ejecta interacting with non-uniform ambient ISM \\citep{braun86}. We resolve the complex X-ray structure of G292.0+1.8: i.e., metal-rich ejecta, shocked dense asymmetric CSM, and the shocked stellar winds or ISM with the high resolution ACIS data. We should note, however, that EW depends not only on the abundance, but also on the temperature and ionization state, so that direct spectral analysis is needed to confirm the inferred metallicity distributions by the EW maps. Further progress on the origin of the equatorial belt and the soft shell-like structure at the SNR's periphery will benefit from follow-up optical spectroscopy and detailed X-ray spectral analysis (H01b)." + }, + "0112/astro-ph0112205_arXiv.txt": { + "abstract": "Archeops, a balloon-borne experiment, will provide a measurement of CMB anistropies from large to small angular scale thanks to its large sky coverage (30\\%), its high angular resolution (10 arcminutes), and its high signal-to-noise ratio due to high sensitivity 100 mK cooled bolometers. We will therefore be able to put strong constraints on the value of the cosmological parameters. Archeops flew already twice, once in Sicily for a technical flight in 1999, and once from Sweden for its first scientific flight in January 2001. I describe here Archeops'main characteristics, the preliminary results from the scientific flight, the expected precision of this flight for power spectrum measurements, and perspectives for the next flights this winter. ", + "introduction": "The Cosmic Microwave Background (CMB) was discovered by Penzias and Wilson \\cite{PenziasWilson:1965} in 1965, and interpreted by Dicke et al. \\cite{Dickeetal:1965}. This radiation comes from the first moments of our Universe, and was first predicted by Gamow, Alpher and Hermann \\cite{AlpherHermGamow:1948} in 1948, in the context of the Big Bang theory. The Big Bang theory predicted that our Universe was in expansion, which was first observed by Hubble \\cite{hubble:1929} in 1929, via the redshift of galaxies. As a consequence of the expansion, the Universe also cools down, meaning that at earlier times it was hotter. Going back in the past, our Universe was so hot that matter and photons were tightly coupled with each other. They formed a plasma in thermal equilibrium. As the Universe cooled down, the photon energy became too small to ionise the matter, below 0.1 eV (below 13.6 eV due to the high photon to electron ratio), the mean free path of the photons became larger than the horizon. The photons were free to cross the Universe towards our detectors. The blackbody dsitribution and spatial properties of these photons remain unchanged due to their negligeable cross-section with matter, the blackbody temperature is just cooled to 2.7 K due to the Universe expansion. Through this radiation, we obtain a picture of our Universe 300 000 years after the Big Bang. ", + "conclusions": "Archeops is a balloon-borne experiment, designed to measure the CMB anistropies from large (about 10 degrees) to small angular scales (about 10 arcminutes). Archeops is also a testbed for Planck HFI, and will provide useful information on galactic dust and point source polarisation. The balloon has already flown twice successfully: once during a test flight and another time for its first scientific flight from Kiruna. The first scientific flight yielded 7.5 hours of scientific data, with 14 detectors dedicated to CMB measurements. Results from this flight will come soon, and we showed that it will give good constrains on large angular scale, but less constraining on small angular scale due to the short duration of the flight. Archeops will fly again this Winter, in order to increase the accuracy on large scales and measure the small scales." + }, + "0112/astro-ph0112175_arXiv.txt": { + "abstract": "We predict from a survey of equations of state and observations of X--ray pulsations from SAX J1808.4--3658, that the upper limit of the mass of the compact star is 2.27 \\msun. The corresponding upper limit of the radius comes out to be 9.73 km. We also do a study to estimate the lower limit of the mass of the compact star. Such a limit is very useful to put constraints on equations of state. We also discuss the implications of the upper mass limit for the detection of the source in radio frequencies. We point out that the possible observation of radio--eclipse may be able to rule out several soft equation of state models, by setting a moderately high value for the lower limit of inclination angle. ", + "introduction": "\\label{sec: 1} The discovery of millisecond X--ray pulsations (period $T = 2.49$ ms; Wijnands \\& van der Klis 1998) in the transient X--ray burster SAX J1808.4--3658 confirmed the speculation that LMXBs are progenitors of millisecond pulsars (Bhattacharya \\& van den Heuvel 1991). The orbital period ($P_{\\rm orb} = 2.01$ hr) and the pulsar mass function ($f_{\\rm 1} = 3.7789 \\times 10^{-5}$) of this source were observationally determined by Chakrabarty \\& Morgan (1998). These give valuable information about the masses (of both the primary and the secondary) and the inclination angle. For example, the value of $P_{\\rm orb}$ uniquely determines the mass of a Roche lobe filling low mass star with known mass--radius relation. It has been recently proposed that the compact star in SAX J1808.4--3658 is a strange star (SS) and not a neutron star (NS) (Li et al. 1999). Such a speculation, if confirmed, will prove that the so called {\\it strange matter hypothesis} (Witten 1984) is correct. According to this hypothesis, strange quark matter (made entirely of deconfined {\\it u, d} and {\\it s} quark) could be the true ground state of strongly interacting matter rather than $^{56}$Fe. This is an important problem of the fundamental physics. To resolve it, we need to constrain the equations of state (EOS) for this compact star very effectively. In this paper, we estimate the upper limits of the mass and the radius of the compact star in SAX J1808.4--3658. We also discuss the possible ways to estimate lower mass limit. ", + "conclusions": "\\label{sec: 4} In this paper, we have estimated the upper limits of the mass and the radius of the compact star in SAX J1808.4-3658. Li et al. (1999) have concluded that a narrow region in $m_{\\rm 1}-R_{\\rm 1}$ space will be allowed for this star. The upper boundary of the mass will constrain this region effectively. It can also give the upper limit of $i$ (from eqn. 3), if $m_{\\rm 2, min}$ is known by an independent measurement. Alternatively, $m_{\\rm 1, max}$ gives the upper limit of $m_{\\rm 2}$, for a known value of $i_{\\rm min}$. For example, $i_{\\rm min} = 63^o$ gives $m_{\\rm 2, max} = 0.066$ for $m_{\\rm 1, max} = 2.27$. It has been proposed that SAX J1808.4-3658 may emerge as a radio pulsar during the X--ray quiescence (Chakrabarty \\& Morgan 1998). Ergma \\& Antipova (1999) have calculated that for $\\lambda < 3$ cm, it may be possible to observe radio emission from this source. However our limits of mass values give a slightly higher $(3.8$ cm) upper limit for $\\lambda$. We have already mentioned in the previous section that a moderately high value of $i_{\\rm min}$ will give a lower limit of $m_{\\rm 1}$. This will be very important for constraining EOS more decisively (as corresponding to every EOS, there exists a maximum possible mass). For example, if $i_{\\rm min} = 63^o$ (corresponds to $m_{\\rm 1, min} = 1.48$, given in the previous section), our EOS models SS1, SS2 and Y will be unfavoured (see the `$m_{\\rm 1, max}$' column of Table 1). According to Chakrabarty \\& Morgan (1998), a deeper eclipse might be observed for the less penetrating radio emission, providing a strong constraint on the value of $i$. Therefore, we expect that, the value of $i_{\\rm min}$ (determined by this method) may be able to rule out several soft EOS models in future." + }, + "0112/astro-ph0112496_arXiv.txt": { + "abstract": "Broadband optical spectra are presented for 34 known and candidate nearby stars in the southern sky. Spectral types are determined using a new method that compares the entire spectrum with spectra of more than 100 standard stars. We estimate distances to 13 candidate nearby stars using our spectra and new or published photometry. Six of these stars are probably within 25~pc, and two are likely to be within the RECONS horizon of 10 pc. ", + "introduction": "The nearest stars have received renewed scrutiny because of their importance to fundamental astrophysics ({\\it e.g.}, stellar atmospheres, the mass content of the Galaxy) and because of their potential for harboring planetary systems and life ({\\it e.g.}, the NASA Origins and Astrobiology initiatives). The smallest stars, the M dwarfs, account for at least 70\\% of all stars in the solar neighborhood and make up nearly half of the Galaxy's total stellar mass (Henry {\\it et al.} 1997; hereafter H97). Their slightly lesser cousins, the brown dwarfs, may lurk in comparable numbers. Yet, many of the nearest red, brown, and white dwarfs remain unrecognized because of their low luminosities. H97 estimate that more than 30\\% of stellar systems within 10~pc of the Sun are currently missing from compendia of nearby stars. The number of ``missing'' stars within 25~pc of the Sun is estimated to be twice the fraction missing within 10~pc. The NASA/NSF NStars Project is a new effort to foster research on all stars within 25~pc, with special emphasis on the development of a comprehensive NStars Database. All systems with trigonometric parallaxes greater than or equal to 0\\arcpt04000 from {\\it The Yale Catalog of Stellar Parallaxes} (YPC, van Altena {\\it et al.} 1995) and {\\it The Hipparcos Catalogue} (HIP, ESA 1997) have been included in the Database. The weighted means of the YPC and HIP parallaxes have been determined, including the combination of all trigonometric parallax values for stellar systems in which widely separated components have had separate parallax measurements. Table~1 lists the numbers of known and predicted stellar systems within 25~pc, and their distributions within equal regions of the sky, obtained from the NStars Database as of 01 July 2001. The predicted number of 1375 systems in each region is based on the assumptions that (1) the density of stellar systems within 5~pc (0.084 systems~pc$^{-3}$) extends to 25~pc, and (2) the distribution of the systems is isotropic. Table~1 clearly shows that more stars are missing in the southern sky than in the northern sky --- we predict that more than two-thirds of the systems are undiscovered in the south. Furthermore, new systems within 5~pc are still being found (H97), so the total number predicted within 25~pc is a lower limit. In a concerted effort to discover and characterize the nearest stars, the RECONS (Research Consortium on Nearby Stars) team has been conducting astrometric, photometric, spectroscopic, and multiplicity surveys of known and candidate stars within 10~pc (for more information about RECONS, see H97). In this paper, the sixth in {\\it The Solar Neighborhood} series, we present optical spectra of 34 known or suspected nearby southern red and white dwarfs, including 10 known members of the RECONS sample and 16 stars for which no spectral types have been previously published. We report spectral types for all the stars in our sample using a method that will define the spectral types used in the NStars Database. We supplement the spectral data with VRI photometry for five stars. Our analysis has revealed two new stars that are probably closer than the 10~pc RECONS horizon, and four others that are probably closer than the 25~pc NStars horizon. ", + "conclusions": "The recent identifications of candidate nearby stars from the proper motion studies of Wroblewski {\\it et al.}, Ruiz {\\it et al.}, and others, and from photometric sky surveys such as DENIS, 2MASS and SDSS, suggest that many nearby stars remain undiscovered. In essence, this paper represents a small step in fingerprinting some suspected nearby stars via spectroscopy. We have established a method for consistent spectral typing that will provide definitive types both for the RECONS effort (horizon 10~pc) and the more extensive NASA/NSF NStars Project (horizon 25~pc). Using this method, we report the first spectral types on a standard system for 16 nearby star candidates. We also provide updated spectral types for 18 other stars using broader spectral coverage than was previously available. This work will allow us to improve the luminosities, colors, and temperatures for the ubiquitous red dwarfs, as well as broaden the database used to investigate the luminosity function, mass function, kinematics, and multiplicity of stars in the solar neighborhood. The nearest objects, such as GJ~1123 and GJ~1128 from this study, will be prime targets of upcoming NASA missions like SIRTF, SIM, and TPF, as well as being additions to the target lists of SETI efforts like Project Phoenix." + }, + "0112/astro-ph0112343_arXiv.txt": { + "abstract": "We present the first results of a deep WFPC2 photometric survey of the loose galactic globular cluster NGC 288. The fraction of binary systems is estimated from the color distribution of objects near the Main Sequence (MS) with a method analogous to that introduced by \\cite{ruba1}. We have unequivocally detected a significant population of binary systems which has a radial distribution that has been significantly influenced by mass segregation. In the inner region of the cluster ($r<1 r_h \\simeq 1.6 r_c$) the binary fraction ($f_b$) lies in the range 0.08--0.38 regardless of the assumed distribution of mass ratios, $F(q)$. The most probable $f_b$ lies between 0.10 and 0.20 depending on the adopted $F(q)$. On the other hand, in the outer region ($r\\ge 1 r_h $), $f_b$ must be less than 0.10, and the most likely value is 0.0, independently of the adopted $F(q)$. The detected population of binaries is dominated by primordial systems. The specific frequency of Blue Straggler Stars (BSS) is exceptionally high, suggesting that the BSS production mechanism via binary evolution can be very efficient. A large population of BSS is possible even in low density environments if a sufficient reservoir of primordial binaries is available. The observed distribution of BSS in the Color Magnitude Diagram is not compatible with a rate of BSS production which has been constant in time, if it is assumed that all the BSS are formed by the merging of two stars. ", + "introduction": "Binary systems are the most frequent form in which stars present themselves, at least in our local neighborhood \\citep{duqma}, and it is generally believed that the same process of stellar birth commonly takes place in clusters \\cite[e.g.,][and references therein]{math}. Furthermore, binaries have a significant impact on the chemical evolution of galaxies and even toy chemical evolution models can easily show that our Universe would not have been as it is without them \\citep{port00}. In collisional stellar systems, binaries can play a key role also in the dynamical evolution. In particular, binaries provide the gravitational fuel that stops and eventually reverses the process of core collapse in globular clusters \\cite[][and references therein]{hual,mh97}. Moreover, the evolution of binaries in clusters can produce peculiar stellar species such as Blue Stragglers (BSS), Cataclysmic Variables (CV), Low Mass X ray sources (LMXB), millisecond pulsars (MSP) and possibly sdB \\citep{bai95,port97a,port97b,sgb,maxt01,gre01,gls}. Thus, the study of binary populations in globular clusters can provide powerful constraints both on dynamical models or on models of formation of exotic objects. Despite their potential interest, the actual detection of binaries and estimate of the binary fraction ($f_B$)\\footnote{Defined as the ratio between the number of binary systems and the total number of cluster members (i.e. binary systems + single stars), see \\cite{hual}.} in globular clusters has eluded the effort of researchers until very recent times \\cite[see][]{hual,gg,itr}, because of the challenging observational requirements. There are three main techniques which have been used to detect binary populations in globular clusters (reviewed by \\citep{hual}): (i) radial velocity variability surveys \\citep{lath}, (ii) searchs for eclipsing variables \\citep{mat96a}, and (iii) searches for a Secondary Main Sequence \\cite[SMS, see][]{rw91,tout99} parallel to the normal Main Sequence (MS) in the Color Magnitude Diagram (CMD). Methods (i) and (ii) are based on the actual detection of individual binary systems. They require large amounts of observing time, since time series are necessary. Method (i) also requires very high precision radial velocity measures that limit the luminosity range of the targets and place great demands on the quality and stability of the instrumental setup. These lead to observational biases and intrinsic limitations \\cite[see \\S2.1 and \\S2.2 in][]{hual} and ultimately to a low discovery efficiency. Thus, while these methods provide the only route to the physical characterization of individual systems (masses, orbits etc.), they are maladapted to determining population properties, such as the binary fraction. In contrast, method (iii) is statistical in nature and does not need repeated observations. It is based on the simple fact that any binary system at the distance of globular clusters is seen as a single star with a flux equal to the sum of the fluxes of the two components. In the CMD, systems composed of two equal mass MS stars lie on a sequence $0.752$ mag brighter than the single stars sequence. If the masses are not equal the shift away from the MS is smaller and depends non-linearly on the mass ratio of the two components, $q={M_2\\over{M_1}}\\le 1.0$ \\cite[see][]{tout99}, where $M_1$ and $M_2$ are the mass of the primary and secondary stars, respectively. While the SMS can detect binaries of any orbital period and any orientation of the orbital plane, it suffers from three problems: \\begin{itemize} \\item The observational signatures of a genuine binary system and two blended, but unrelated, single stars are indistinguishable. Thus, any SMS in a globular cluster is contaminated to some degree by blended objects. \\item Exceptionally accurate photometry extending to at least a couple of magnitudes below the Turn Off point (TO) is required. Any secondary sequence can be easily obscured by the observational scatter in the single star sequence. Thus, while hints of a SMS have been found in many cases \\cite[see Tab. 3 in][]{hual}, the only definitive detections from ground based photometry have been in two very loose and relatively nearby clusters, E3 \\citep{ver96,mcc85} and NGC 288 \\citep{bol92}. \\item The distribution of mass ratios must be retained as a free parameter while estimating $f_b$, although some observational constraints on this distribution may emerge. \\item SMS is very sensitive to the distribution of the mass ratios. To obtain an useful constraint on the binary population, the high $q$ part of the distribution must be significantly populated. \\end{itemize} Sufficient photometric accuracy at such faint magnitudes can often be be achieved with HST observations, and sometimes with large ground based telescopes in optimal seeing conditions. The correction for blendings has a very complex behavior and can be made only via extensive artificial star experiments which precisely mimic the observations and data reduction. While the above framework was clearly recognized, e.g., by \\cite{bol92}, the full development and application of a method that properly accounted for all these effects emerged only with the seminal work by \\cite{ruba1}. These authors analyzed a large set of exposures taken with the Planetary Camera of the HST-WFPC2, imaging a field including the central region of the globular cluster NGC~6752. They were able to estimate $f_b$ from the distribution of deviation in color with respect to the MS ridge line using a large set of artificial stars experiments to correct for blendings and considering different possible distributions of the mass ratios. They found 15\\% $\\le f_b \\le 38$ \\% for the sample within the core radius of NGC~6752 and marginal evidence for $ f_b \\le 16$ \\% outside that limit. We have used a method very similar to that of \\cite{ruba1} [RB97] on HST observations of two partially overlapping WFPC2 fields sampling the central part of the low density cluster NGC288. In \\S2 we will describe the observations, data reduction, and artificial star experiment. In \\S3 the method will be described in detail, outlining the differences with respect to RB97. Results and discussion are reported in \\S4. Once the binary fraction and its spatial properties are established, we turn to the analysis of the expected products of the evolution of binary systems, in particular BSS (\\S5). Finally we summarize the results and provide a global description of the status and the evolution of the binary population in NGC288 (\\S6). Some preliminary results from the observations presented in this paper have been reported in a recent meeting \\citep{bm00}. ", + "conclusions": "We have measured the fraction of binary system ($f_b$) in the loose globular cluster NGC~288 using the SMS method to identify binaries in the $V,~V-I$ CMD. We employed a technique that accurately accounts for {\\em all} the observational effects (observational scatter, blending, etc., see RB97) that may affect the estimate of $f_b$. This is only the second measurement of this kind ever obtained for a globular cluster, the first having been made by RB97 for NGC~6752. We find that the observations are strongly incompatible with the hypothesis $f_b = 0$ \\%, independently of any assumption concerning the distribution of mass ratios, $F(q)$, thus {\\em the presence of binary systems in NGC 288 is confirmed beyond any doubt} \\cite[see also][]{bol92}. We have estimated $f_b$ for various assumed $F(q)$, and have found that, independently of the adopted $F(q)$, {\\em the observations are strongly incompatible with global binary fractions lower than 5--7\\% or higher than $\\sim 30$ \\%}. The binary fraction in the region within the half-light radius ($r_h$) is significantly larger than in the region outside this limit. For $r < 1\\, r_h$, 8 \\% $< f_b \\le 38$ \\% and the most probable binary fractions ranges from 10 \\% to 25 \\%, depending somewhat on the assumed $F(q)$. On the other hand, for $r > 1\\, r_h$, $f_b \\le 10$\\% and the most probable binary fraction is $f_b = 0$ \\%, independently of the assumed $F(q)$. Hence, {\\em binary systems are much more abundant in the inner regions of the cluster, a clear sign of the occurrence of dynamical mass segregation}. Simple dynamical arguments strongly suggest that {\\em the large majority of binary systems present in NGC~288 are of primordial origin}, and that single star - single star collision processes are highly inefficient in this cluster. Thus the large majority of the BSS in NGC~288 must have a binary origin. Despite that, {NGC~288 has a very high Specific Frequency of BSS, comparable to or exceeding that of much more dense clusters}. Like the binaries, the BSS population is centrally concentrated with virtually all the identified BSS lie within $1\\,r_h$ from the cluster center. The selected BSS candidates appear to have masses between $\\sim 0.9 M_{\\odot}$ and $\\sim 1.4 M_{\\odot}$, and form a remarkably narrow and well defined sequence in the $(V,~V-I)$ CMD. If the majority of BSS have been produced by the merging of two binary member stars, then the {\\em BSS distribution is not compatible with a BSS formation rate which has been constant in time}. On the contrary, the existence of a significant peak in the BSS formation rate in the recent past ($\\sim 1$--4 Gyr ago and lasting $\\sim 1$--2 Gyr) is suggested by the comparison with theoretical models \\citep{sills00}. The observations also seem compatible with a scenario in which most of the BSS population of NGC~288 has been producted via the mass transfer occurring in close binary systems. A few {\\em yellow stragglers} and one candidate E-BSS (in the Helium burning phase) have also been identified. \\subsection{Binary Fraction: Comparisons with Other Systems} The two globular clusters for which a robust estimates of the {\\em global} binary fraction have been obtained using the SMS technique, i.e., NGC~6752 (RB97) and NGC~288 (this work) are very different in terms of stellar density. The central density in NGC~6752 is $\\sim 1000$ times that of NGC~288. It is very likely that the past dynamical evolution of the two clusters has been very different. Despite of that, the present day observed binary fraction (and binary distribution) is remarkably similar, at least at the present level of accuracy. However, while many details of the radial distribution of the binary populations are beyond the reach of the present day techniques, a first comparison of their broad properties is now possible. For NGC~6752 $15 \\% \\le f_b \\le 38 \\%$ within $1\\, r_c$ from the cluster center and $f_b \\le 16 \\%$ in the outer region. For NGC~288 8 \\% $\\le f_b \\le 38$ \\% within $1\\, r_h \\simeq 1.6\\, r_c $ and $f_b \\le 10$ \\% in the outer region. It may be presumed that a broad upper limit is set by primordial conditions and by the fact that in any globular cluster a significant number of binaries are {\\em soft} (and thus are rapidly destroyed). Lower limits may be set by the equilibrium between destruction and formation of binary systems, somehow regulated by the density of the environment \\cite[since the efficiency of both mechanisms increase with stellar density,] []{hual}. It is also interesting to note that the global $f_b$ extrapolated from spectroscopic binaries and/or eclipsing variable searches in globulars is broadly constrained to be in the range 10 \\% $ \\le f_b \\le 40 $ \\% \\cite[see][and references therein]{mat96a}. Hence, two decades after the pioneering (and negative) results by \\citet{gg}, a very different picture of binaries in globular clusters seems firmly established: {\\em the binary fraction in globulars is not null (and likely larger than $\\sim 5-10$ \\%), but is still significantly lower than in most of other environments}, i.e., the local field, open clusters and star forming regions, for which $f_b\\ge50$ \\% \\citep{mat96a,math,duqma}. \\subsection{Blue Stragglers in Low Density Systems} One of the most interesting results of the present analysis is the demonstration that {\\em the formation of BSS via mass transfer/coalescence of primordial binary systems may be as efficient as collisional mechanisms, occurring in the most favorable environments} (see \\S5). This is a further (and quantitative) piece of evidence indicating that {\\em large populations of BSS may be produced in environments with remarkably low stellar density, if a sufficient reservoir of primordial binaries is available} \\citep[see also][and references therein]{nh87,nc89,leolin,leo93,mat96a,mat96b}. This statement may have important consequences in the interpretation of the ubiquitous {``blue plumes''} observed above the main MSTO in the CMD of many dwarf Spheroidal galaxies (dSphs) dominated by old ``globular-cluster-like'' populations \\cite[as for instance Ursa Minor, Draco, Sextans and Sagittarius, see][respectively, and references therein]{mart01,cs86,mat91,sgr2}. These sparsely populated sequences are usually associated with recent, very small episodes of star formation. This hypothesis is also supported by the presence of stellar species thought to result from the evolution of stars with initial mass larger than the typical TO mass of old population (1--$2\\,M_{\\odot}$ vs. $0.8\\, M_{\\odot}$), such as Anomalous Cepheids (AC) and/or bright AGB stars \\cite[see][and references therein]{mat98,dacosta}. Even though the blue plumes lie in the region of the CMD populated by BSS, BSS are sometimes considered an unlikely explanation because of the very low density environments. The results presented here confirms that a dense environment is not a {\\em conditio sine qua non} for the efficient production of BSS. Further, there is no reason to believe that primordial binaries are under-abundant in dSphs \\citep[see also][and references therein]{leof,leo93}. Indeed, the few available estimates suggest that (at least in Ursa Minor and Draco) the binary fraction may be even larger than in the local field \\citep{opa}. It is also important to recall that once a star more massive than the typical old TO stars is formed from a binary system (via mass transfer or coalescence), it will follow the evolutionary path typical of its {\\em new} mass, thus possibly passing through the AC or bright-AGB phases \\cite[as succesfully demonstrated by][more than two decades ago]{rms}. There is every reason to suspect a significant BSS population in dSphs, and they remain a viable explanation for the blue plumes \\citep[see, e.g.][]{grill}." + }, + "0112/astro-ph0112357_arXiv.txt": { + "abstract": "I review what we currently do and do not know about the masses of disk galaxies and their dark matter halos. The prognosis for disks is good: the asymptotic rotation velocity provides a good indicator of total disk mass. The prognosis for halos is bad: cuspy halos provide a poor description of the data, and the total mass of individual dark matter halos remains ill-constrained. ", + "introduction": "The great regularity of the Tully-Fisher relation \\cite{TF} has long been thought to originate from a strong mass-velocity relation and a near constancy of mass-to-light ratio. The latter requires a fair but not unreasonable amount of regularity in stellar populations. Put simply,\\footnote{``Stacy, you're a genius! ...[!] ...when it comes to pepper grinders'' (van den Bosch 2001, private communication).} \\begin{equation} L \\sim {\\cal M} \\sim V^a \\; . \\end{equation} There have long been indications (Sancisi 1995, private communication) that this simple scaling may fail at low luminosities. This has become more clear as data have improved \\cite{MG},\\cite{Stil}. This breakdown of the Tully-Fisher relation might arise because of the chaotic star formation histories of low mass galaxies, or as a result of a breakdown in the underlying mass-velocity relation. Another possibility is that optical luminosity ceases to trace mass because stars cease to be the dominant mass component in these disks \\cite{Ken}. It has now become clear that this last possibility is in fact the case. Low mass galaxies are often dominated by gas rather than stars. If instead of luminosity or stellar mass, we plot disk (star + gas) mass against the flat rotation velocity, a nice mass-velocity relation is recovered over many orders of magnitude (Fig.~1). This `Baryonic Tully-Fisher Relation' (BTF) is \\cite{MSBB} \\begin{equation} {\\cal M}_d = {\\cal A} V_f^b \\; , \\end{equation} for which the data in Fig.~1 give \\begin{subeqnarray} {\\cal A} = 50\\;{\\cal M}_{\\odot}\\;{\\rm km}^{-4}\\;{\\rm s}^4 \\nonumber \\\\ b = 4.0 \\pm 0.1 \\; . \\setcounter{eqsubcnt}{0} \\end{subeqnarray} \\begin{figure}[t] \\begin{center} \\includegraphics[width=.9\\textwidth]{mcgaughF1.eps} \\end{center} \\caption[]{The Tully-Fisher relation expressed in terms of (\\textbf{a}) stellar mass and (\\textbf{b}) baryonic disk mass (the sum of stars and gas). The luminous Tully-Fisher relation holds well for galaxies dominated by stars, but breaks down for low mass galaxies where the gas mass can often exceed the stellar mass (\\textbf{a}). The sum of stars and gas provides a better correlation (\\textbf{b}): the asymptotic flat rotation velocity is a good indicator of disk mass \\cite{MSBB}. The data shown here are taken from a large compilation of high quality data \\cite{SM}. Consequently, the scatter is greatly reduced from that in \\cite{MSBB}. The galaxies shown here are drawn from the full range of disk Hubble types: mostly Sc, Sd, Sm, Im, but also a few Sa and Sb galaxies are present. The intrinsic scatter is small, with room only for scatter in the stellar mass-to-light ratio due to variations in star formation histories (probably not in the IMF), and scatter due to the modest ellipticities of disks \\cite{Bersh}. } \\label{fig1} \\end{figure} The normalization of the BTF is rather uncertain: formally acceptable values fall in the range $34 < {\\cal A} < 85$. The precise value of the slope has been modestly controversial: $b= 4.0$ was given by \\cite{MSBB} while $b=3.5$ was found by \\cite{BJ}. This difference can be traced to different assumptions about the (rather goofy \\cite{HSTKP}) distance to the UMa cluster for which some of the better rotation curve \\cite{VS} and photometric data \\cite{TVPHW} exist. As the distance increases, the gas mass increases faster than the stellar mass (as $D^2$ and as $D$, respectively). This boosts the total mass of gas dominated galaxies by a larger factor than star dominated galaxies. Since these reside at opposite ends of the relation, the slope tips to shallower values with increasing $D$. Nevertheless, the population models of \\cite{BJ} are consistent with a slope of $b= 4.0$ (Fig.~2). While the calibration of the BTF can always be improved, it already provides an excellent indicator of disk mass. Moreover, continuity between gas-rich and star-rich galaxies constrains stellar population mass-to-light ratios. The favored values are reasonable in terms of population synthesis models (Fig.~2), but unpleasantly heavy for cuspy dark matter halos. \\begin{figure}[t] \\begin{center} \\includegraphics[width=.9\\textwidth]{mcgaughF2.eps} \\end{center} \\caption[]{The stellar mass-to-light ratios in (\\textbf{a}) the $B$-band and (\\textbf{b}) the $K'$-band predicted by a slope 4 BTF for the UMa galaxies \\cite{Verh}, \\cite{TVPHW}, \\cite{VS}. These are plotted as a function of $B-V$ color, together with the Bruzual \\& Charlot, Salpeter IMF model from \\cite{BJ} (the first model in their table 4). The population synthesis models are in good agreement with the BTF, indicating that we have a good handle on ${\\cal M}_*/L$ and disk masses.} \\label{fig2} \\end{figure} ", + "conclusions": "" + }, + "0112/astro-ph0112482_arXiv.txt": { + "abstract": "We describe an automated method, the Cut \\& Enhance method (CE) for detecting clusters of galaxies in multi-color optical imaging surveys. This method uses simple color cuts, combined with a density enhancement algorithm, to up--weight pairs of galaxies that are close in both angular separation and color. The method is semi--parametric since it uses minimal assumptions about cluster properties in order to minimize possible biases. No assumptions are made about the shape of clusters, their radial profile or their luminosity function. The method is successful in finding systems ranging from poor to rich clusters of galaxies, of both regular and irregular shape. We determine the selection function of the CE method via extensive Monte Carlo simulations which use both the real, observed background of galaxies and a randomized background of galaxies. We use position shuffled and color shuffled data to perform the false positive test. We have also visually checked all the clusters detected by the CE method. We apply the CE method to the 350 deg$^2$ of the SDSS (Sloan Digital Sky Survey) commissioning data and construct a SDSS CE galaxy cluster catalog with an estimated redshift and richness for each cluster. The CE method is compared with other cluster selection methods used on SDSS data such as the Matched Filter (Postman et al.\\ 1996, Kim et al.\\ 2001) , maxBCG technique (Annis et al.\\ 2001) and Voronoi Tessellation (Kim et al.\\ 2001). The CE method can be adopted for cluster selection in any multi-color imaging surveys. ", + "introduction": "\\label{Jun 28 09:03:18 2001} Clusters of galaxies are % the most massive virialized systems known and provide powerful tools in the study of cosmology and extragalactic astronomy. For example, clusters are efficient tracers of the large--scale structure in the Universe as well as determining the amount of dark matter on Mpc scales (Bahcall 1998; Carlberg et al. 1996; Borgani \\& Guzzo 2000 and Nichol 2001 and references therein). Furthermore, clusters provide a laboratory within which to study a large number of galaxies at the same redshift and thus assess the effects of dense environments on galaxy evolution {\\it e.g.} morphology--density relation (Dressler et al.\\ 1980, 1984, 1997), Butcher--Oemler effect (Butcher \\& Oemler 1978, 1984) and the density dependence of the luminosity function of galaxies (Garilli et al. 1999). In recent years, surveys of clusters of galaxies have been used extensively in constraining cosmological parameters such as $\\Omega_m$, the mass density parameter of the universe, and $\\sigma_8$, the amplitude of mass fluctuations at a scale of 8 $h^{-1}$ Mpc (see Oukbir \\& Blanchard 1992; Viana \\& Liddle 1996, 1999; Eke et al. 1996; Bahcall, Fan \\& Cen 1997; Henry 1997, 2000; Reichart et al. 1999 as examples of an extensive literature on this subject). Such constraints are achieved through the comparison of the evolution of the mass function of galaxy clusters, as predicted by the Press-Schechter formalism (see Jenkins et al. 2001 for the latest analytical predictions) or simulations ({\\it e.g.} Evrard et al. 2001 and Bode et al. 2001), with the observed abundance of clusters with redshift. Therefore, to obtain robust constraints on $\\Omega_m$ and $\\sigma_8$, we need large samples of clusters that span a large range in redshift and mass as well as possessing a well--determined selection function (see Nichol 2001). Despite their importance, existing catalogs of clusters are limited in both their size and quality. For example, the Abell catalog of rich clusters (Abell 1958), and its southern extension (Abell, Corwin and Olowin 1989), are still some of the most commonly used catalogs in astronomical research even though they were constructed by visual inspection of photographic plates. Another large cluster catalog by Zwicky et al.\\ (1961-1968) was similarly constructed by visual inspection. Although the human eye can be efficient in detecting galaxy clusters, it suffers from subjectivity and incompleteness. For cosmological studies, the major disadvantage of visually constructed catalogs is the difficulty to quantify selection bias and thus, the selection function. % Furthermore, the response of photographic plates is not uniform. Plate-to-plate sensitivity variations can disturb the uniformity of the catalog. To overcome these problems, several cluster catalogs have been constructed using automated detection methods on CCD imaging data. % They have been, however, restricted to small areas due to the lack of large--format CCDs. {\\it e.g.} the PDCS catalog (Postman et al.\\ 1996) only covers 5.1 deg$^2$ with 79 galaxy clusters. % The need for a uniform, large cluster catalog is strong. The Sloan Digital Sky Survey (SDSS; York et al.\\ 2000) data offer the opportunity to produce the largest and most uniform galaxy cluster catalog in existence because the SDSS is the largest CCD imaging survey currently underway scanning 10,000 deg$^2$ centered approximately on the North Galactic Pole. The quantity and quality of the SDSS data demands the use of sophisticated cluster finding algorithms to help maximize the number of true cluster detections while suppressing the number of false positives. The history of the automated cluster finding methods goes back to Shectman's count-in-cell method (1985). He counted the number of galaxies in cells on the sky to estimate the galaxy density. Although this provided important progress over the visual inspection, the results depend on the size and position of the cell. Currently the commonly used automated cluster finding method is the Matched Filter technique (MF) (Postman et al.\\ 1996, Kawasaki et al.\\ 1998, Kepner et al.\\ 1999, Schuecker \\& Bohringer 1998, Bramel et al.\\ 2000, Lobo et al.\\ 2000, da Costa et al.\\ 2000 and Willick et al.\\ 2000). The method assumes a filter for the radial profile of galaxy clusters and for the luminosity function of their members. It then selects clusters from imaging data by maximizing the likelihood of matching the data to the cluster model. Although the method has been successful, galaxy clusters that do not fit the model assumption (density profile and LF) may be missed. We present here a new cluster finding method called the {\\it Cut } and {\\it Enhancement } method, or CE. This new algorithm is semi--parametric and is designed to be as simple as possible using the minimum number of assumptions possible about cluster properties. In this way, it should be sensitive to all types of galaxy overdensities even those that may have recent under--gone a merger and therefore, are highly non--spherical. One major difference between CE and previous cluster finders is that CE makes full use of colors of galaxies, which become available due to the advent of the accurate CCD photometry of the SDSS data. We apply this detection method % on 350 deg$^2$ of the SDSS commissioning data and construct the large cluster catalog. The catalog ranges from rich clusters to the more numerous poor clusters of galaxies over this area. We also determine the selection function of the CE method. In Sect.\\ \\ref{sec:data}, we describe the SDSS commissioning data. In Sect.\\ \\ref{sec:method}, we describe the detection strategy of Cut \\& Enhance method. In Sect.\\ \\ref{sec:monte}, we present the performance test of the Cut \\& Enhance method and selection function using Monte Carlo simulations. In Sect.\\ \\ref{sec:visual}, we visually check the success rate of the Cut \\& Enhance method. In Sect.\\ \\ref{sec:compare}, we compare Cut \\& Enhance method with the other detection methods applied to the SDSS data. In Sect.\\ \\ref{sec:summary}, we summarize the results. ", + "conclusions": "\\label{sec:summary} We have developed a new cluster finding method, the Cut \\& Enhance method. It uses 30 color cuts and four color-color cuts to enhance the contrast of galaxy clusters over the background galaxies. After applying the color and color-color cuts, the method uses the color and angular separation weight of galaxy pairs as an enhancement method to increase the signal to noise ratio of galaxy clusters. We use the Source Extractor to detect galaxy clusters from the enhanced maps. The enhancement and detection are performed for every color cut, producing 34 cluster lists, which are then merged into a single cluster catalog. Using the Monte Carlo simulations with real SDSS background as well as shuffled background, the Cut \\& Enhance method is shown to have the ability to detect rich clusters ($Ngal$=100) to $z\\sim0.3$ with $\\sim$80\\% probability. The probability drops sharply at $z$=0.4 due to the flux limit of the SDSS imaging data. The positional accuracy is better than 40'' for all richnesses examined at $z\\leq$0.3. The false positive test shows that over 70\\% of clusters are likely to be real systems for CE richness $>$10. We apply Cut \\& Enhance method to the SDSS commissioning data and produce an SDSS Cut \\& Enhance cluster catalog containing 4638 galaxy clusters in $\\sim$350 deg$^2$. We compare the CE clusters with other cluster detection methods: MF, maxBCG and VTT. The SDSS Cut \\& Enhance cluster catalog developed in this work is a useful tool to study both cosmology and property of clusters and cluster galaxies." + }, + "0112/astro-ph0112161_arXiv.txt": { + "abstract": "We compute the bispectrum of the 2dF Galaxy Redshift Survey (2dFGRS) and use it to measure the bias parameter of the galaxies. This parameter quantifies the strength of clustering of the galaxies relative to the mass in the Universe. By analysing 80 million triangle configurations in the wavenumber range $0.1 < k < 0.5\\,h\\,$Mpc$^{-1}$ (i.e. on scales roughly between 5 and 30 $h^{-1}\\,$ Mpc) we find that the linear bias parameter is consistent with unity: $b_1=1.04\\pm 0.11$, and the quadratic (nonlinear) bias is consistent with zero: $b_2=-0.054\\pm 0.08$. Thus, at least on large scales, {\\it optically-selected galaxies do indeed trace the underlying mass distribution}. The bias parameter can be combined with the 2dFGRS measurement of the redshift distortion parameter $\\beta \\simeq \\Omega_m^{0.6}/b_1$, to yield $\\Omega_m = 0.27\\pm 0.06$ for the matter density of the Universe, a result which is determined entirely from this survey, independently of other datasets. Our measurement of the matter density of the Universe should be interpreted as $\\Omega_m$ at the effective redshift of the survey ($z=0.17$). ", + "introduction": "Clustering of mass in the Universe is believed to be a result of amplification by gravitational instability of small perturbations generated in the early Universe. Comparison with theoretical predictions offers the chance to test models of the generation of the perturbations, as well as putting important constraints on cosmological parameters, which control the growth rate of the perturbations. A fundamental limitation on such a comparison has been that theoretical models predict the clustering properties of the mass in the Universe, and yet we have few direct measures of mass observationally. More readily observable is the distribution of luminous objects such as galaxies, so to compare with theory one has to determine, or assume, the relationship between the clustering of mass and the clustering of galaxies. In general one will expect these to differ, because the efficiency of galaxy formation may depend in some nontrivial way on the underlying mass distribution. The idea that structures may be `biased' tracers of the mass distribution goes back to \\scite{Kai84}, who explained the high clustering strength of Abell clusters as due to their forming in high-density regions of the Universe. In addition, observations indicating that different types of galaxy cluster differently (e.g. \\pcite{Dress80,PG84,WTD88,Ham88,LNP90,LahavSaslaw92}) show that they cannot all be unbiased tracers of the mass. Bias became an attractive way to reconcile the low velocities of galaxies with the high-density Einstein-de Sitter model favoured in the 1980s (e.g. \\pcite{DEFW85}), but after the Cosmic Background Explorer (COBE) determined the amplitude of primordial fluctuations on large scales \\cite{COBE}, the `standard' biased Cold Dark Matter (CDM) model became less popular. With the advent of more detailed datasets in the cosmic microwave background (CMB) and large-scale structure, it is possible to investigate and constrain a wider range of galaxy formation models, and an unknown bias relation adds uncertainty to the process. Since the efficiency of galaxy formation is not well understood theoretically, it makes sense to try to measure it empirically from observations. When the perturbations are small (or on large, linear scales), it is difficult to do this: there is a degeneracy between the unknown amplitude of the matter power spectrum $P(k)$ and the degree of bias, $b$, defined such that the galaxy power spectrum is $P_g(k) \\equiv b^2 P(k)$. In principle, $b$ may be a function of scale, through the wavenumber $k$. At later times (or on smaller scales), however, the degeneracy is lifted by nonlinear effects. One feature of nonlinear gravitational evolution is that the overdensity field $\\delta(\\bx) \\equiv (\\rho(\\bx)-\\overline{\\rho})/\\overline{\\rho}\\ $ becomes progressively more skewed towards high density. In principle skewness could also arise from non-Gaussian initial conditions; in practice this can be neglected \\cite{VWHK}, since CMB fluctuations are consistent with Gaussian initial conditions \\cite{Komatsuetal01,Santosetal001}. One can thus hope to exploit the gravitational skewness, but skewness could equally well arise from biasing, e.g. from a galaxy formation efficiency that increased at dense points in the mass field. It is nevertheless possible to distinguish these two effects by considering the {\\it shapes} of isodensity regions. If the field is unbiased, then the shapes of isodensity contours become flattened, as gravitational instability accelerates collapse along the short axis of structures, leading to sheet-like and filamentary structures (e.g. \\pcite{Zel70}). If the galaxy field is highly biased with the same power spectrum, however, the underlying mass field is of low amplitude, and thus expected to be close to the initial field, which is assumed to be Gaussian. These fields do not have highly-flattened isodensity contours, as bias does not flatten the contours; for example Eulerian bias preserves the contour shape. Thus there is a difference which could be detected, for example, by studying the three-point correlation function. In this paper, we exploit this effect in Fourier space rather than real space, by analysing the bispectrum: $\\langle \\delta_{\\bk1}\\delta_{\\bk2} \\delta_{\\bk3} \\rangle$, where $\\delta_\\bk$ is the Fourier transform of the galaxy overdensity field. The theory for the bispectrum is set out in \\scite{Fry94}, \\scite{HBCJ95}, \\scite{MVH97}, \\scite{VHMM98}, \\scite{SCFFHM98}, \\scite{SCF99} and \\scite{Scocc00}. The galaxy survey we use is the Anglo-Australian Telescope 2 degree field Galaxy Redshift Survey \\cite{Colless01}, as compiled in February 2001. It was created with the 2dF multi-fibre spectrograph on the Anglo-Australian Telescope \\cite{Lewisetal2001}, and currently consists of over 200,000 galaxies with redshifts up to about $z=0.3$, broadly in two regions centred near the south and north galactic poles. See {\\tt http://www.mso.anu.edu.au/2dFGRS/} for further details. It is the first survey which is large enough to put tight constraints on the bias parameter, as previous surveys are too shallow or too sparse. In this paper, we use 127,000 galaxies from the February 2001 compilation of the catalogue, truncated at $0.03 1$, where $R$ is the size of an inhomogeneity. Consequently the contribution of higher partial waves to the scattering process cannot be neglected, and gives rise to much larger energy correction, as for large objects more of the incident wave will be reflected \\cite{bma1,bma2,bma3,bma4,bwi}. The leading shell energy contribution, associated with the neutron scattering, to the total energy can be determined in the semiclassical approximation for various geometries of nuclear shapes (see appendix B). It has been shown to yield energy corrections that influence the structure of inhomogeneous nuclear matter \\cite{bma1,bma2,bma3}. It is also obvious that since the gross shell structure is determined by the shortest periodic orbits in the system, the shell energy is a sensitive function of the crystal lattice type, as well as of the shape of inhomogeneities \\cite{bma1,bwi}. Unfortunately the semiclassical approach has drawbacks. First, it assumes that the obstacles are impenetrable scatterers which may overestimate the amplitude of the shell energy. Second, the method lacks of the mutual interplay between the shell energy and the liquid drop energy. Such a coupling is quite important since it allows a part of the shell energy to be ``absorbed'' into deformation. Hence a microscopic treatment of the problem is needed where both effects, i.e. the one coming from the liquid drop energy and the shell energy term are treated on the same footing. ", + "conclusions": "In this paper we have analyzed the structure of the neutron star crust at zero temperature using the Hartree-Fock approach with an effective Skyrme interaction. In particular, we focused on the problem of shell effects associated with unbound neutrons scattered by nuclear inhomogeneities. Our results show that: \\begin{itemize} \\item the relative energies of different phases fluctuate rapidly as a function of the total density, \\item the fluctuations can be attributed to the shell effects associated with unbound neutrons, \\item the neutron shell energy density is of the same order as the liquid drop energy density differences between the various phases present in the crust, and its behavior can be easily understood in terms of the periodic orbits in the system, \\item the pairing correlations lead only to a slight decrease of the shell energy and are not expected to influence significantly the structure of the crust. \\end{itemize} Consequently, we may conclude that the structure of the crust may be quite complicated since the shell effects associated with unbound neutrons may easily reverse the phase transition order predicted by liquid drop based approaches. Moreover the number of phase transitions may increase since the same phase may appear for various density ranges (see Fig. 4). Our results suggest also that several phases, differing by the lattice type, could coexist in the crust. This possibility was not taken into account in previous investigations since in a liquid drop based approach, the system favors only one lattice type for a given nuclear shape. However, since the neutron shell energy is very sensitive to the spatial order in the system, this simple picture will change, and one may expect the coexistence of various lattice geometries with different lattice constants. Note also that once a phase is formed, its geometric order will be stabilized by both the Coulomb repulsion and the shell energy. Indeed, although the Coulomb energy is a smooth function of the nuclear displacement, the shell energy is not. Since several different orbits contribute to the shell effects (except for the slab--like phase) the displacement of a single nucleus from its equilibrium position in the lattice will give rise to interference effects depending on the lattice type. Hence it is likely that the system will favor distorted lattices, or lattices with defects which decrease the shell energy. In that respect the present results support the conclusions drawn on the basis of a simple model \\cite{bma1}. All these arguments suggest that the inner crust is an extremely rich and complicated system which may turn out to be on the verge of a disordered phase." + }, + "0112/astro-ph0112189_arXiv.txt": { + "abstract": "{We use differential CCD photometry to search for variability in $BVI$ among 990 stars projected in and around the old open cluster M\\,67. In a previous paper we reported results for 22 cluster members that are optical counterparts to X-ray sources; this study focuses on the other stars in our observations. A variety of sampling rates were employed, allowing variability on time scales ranging from $\\sim 0.3$ hours to $\\sim 20$ days to be studied. Among the brightest sources studied, detection of variability as small as $\\sigma\\approx$ 10 mmag is achieved (with $>3\\sigma$ confidence); for the typical star observed, sensitivity to variability at levels $\\sigma\\approx 20$ mmag is achieved. The study is unbiased for stars with $12.5 3\\sigma$ confidence) is achieved; for the typical star observed, sensitivity to variability at levels $\\sigma \\approx 20$ mmag is achieved. Membership information is available for 439 stars (46 variables) included in our observations and marks 319 (38 variables) as members with a probability of at least 75\\% (\\cite{san}; \\cite{girard}). Of these 38 variable cluster members, 29 exhibit variability in more than one of the passbands used, increasing our confidence that the observed variability is real. Nine of these stars are periodic variables. In all cases the amplitude of variability is low, ranging from a few hundredths to a few tenths of a magnitude. Our study is sensitive to brightness variations on time scales of 0.3 hours to $\\sim 20$ days. Apparently, at the age of M\\,67 variability on these time scales and at these amplitudes is strongly associated with binarity, as 14 of the 29 ``high-confidence\" variable members are known binaries. One of the other high-confidence variables (\\object{S\\,1112}) is an X-ray source for which binarity has not been established, another is a blue straggler (\\object{S\\,1263}), and still two others are situated on the cluster binary sequence. {\\it Periodic} variability is especially rare for single stars, for in 8 of the 9 periodic variable members observed by us (including the new candidate contact system \\object{S\\,757}), 7 of which are X-ray sources, the variability finds its origin in the binary nature of the stars (eclipses, ellipsoidal variations, rotational spot-modulation in tidally locked binaries). This confirms the picture that rapid rotation in an old population can only be maintained in close binaries. In the contact binary 3665 (\\object{ET\\,Cnc}), for which no membership information exists, the variability is the result of binarity as well. We encourage spectroscopic observations of the three remaining stars that exhibit periodic variations (the faint and blue star 2426 and star 2703 on the binary sequence, both without membership information) and the member 3579 (\\object{S\\,1112}, discussed in paper I) to establish their binary status and/or obtain an indication for membership from their radial velocity. Also, more observations should be obtained of the stars for which we provide tentative periods, in the first place to further examine if their photometric variability is indeed periodic and secondly to establish if they are single or binary. The origin of the photometric variability for the remaining stars discussed in this paper is in most cases unknown. As possible causes for the variations we suggest low-level surface activity, stellar pulsations or, especially for the stars on the binary sequence, binary interaction." + }, + "0112/astro-ph0112140_arXiv.txt": { + "abstract": "The steep spectrum of IRAS\\,F02044+0957 was obtained with the RATAN-600 radio telescope at four frequencies. Optical spectroscopy of the system components, was carried out with the 2.1m telescope of the Guillermo Haro Observatory. Observational data allow us to conclude that this object is a pair of interacting galaxies, a LINER and a HII galaxy, at $z=0.093$. ", + "introduction": " ", + "conclusions": "" + }, + "0112/astro-ph0112376_arXiv.txt": { + "abstract": "The sounds accompanying electrophonic burster meteors are characteristically described as being akin to short duration ``pops'' and staccato--like ``clicks''. As a phenomenon distinct from the enduring electrophonic sounds that occasionally accompany the passage and ablation of large meteoroids in the Earth's lower atmosphere, the bursters have proved stubbornly difficult to explain. A straightforward calculation demonstrates that in contradistinction to the enduring electrophonic sounds, the electrophonic bursters are not generated as a consequence of interactions between the meteoroid ablation plasma and the Earth's geomagnetic field. Here we present a novel and hitherto unrecorded model for the generation of short--duration pulses in an observer's local electrostatic field. Our model is developed according to the generation of a strong electric field across a shock wave propagating in a plasma. In this sense, the electrophonic bursters are associated with the catastrophic disruption of large meteoroids in the Earth's atmosphere. We develop an equation for the description of the electric field strength in terms of the electron temperature and the electron volume density. Also, by linking the electron line density to a meteor's absolute visual magnitude, we obtain a lower limit to the visual magnitude of electrophonic burster meteors of $M_{\\mathrm{v}}\\approx -6.6$, in good agreement with the available observations. ", + "introduction": "While Electrophonic meteor sounds have been widely reported throughout recorded history, they are, none--the--less, a poorly observed phenomena (Keay \\& Ceplecha, \\cite{KEAY3}). By this we mean that the accounts of electrophonic sounds are mostly anecdotal and secondary. To our knowledge only two electrophonic meteors have ever been recorded instrumentally. These are the fireball events of 1981, August 13th as reported by T. Watanabe and co--authors in Japan (see Keay, \\cite{KEAY5} for details) and 1993, August 11th as reported by Beech et al. (\\cite{BEECH1}). Expressed in terms of two broadly divided classes, electrophonic sounds are either of the short duration, or burst type in which a sharp ``click'' of ``pop'' is reported, or of the sustained type in which a temporally extended ``rushing'' or ``crackling'' sound is heard (Keay, \\cite{KEAY4}). For brevity and phenomenological reasons we shall call the short duration electrophonic sounds ``bursters''. The essential characteristics of the electrophonic bursters are their short durations, $\\tau\\approx 1$~s, and their piquant impression on the human auditory system. It is not presently possible to draw any clear statistical inferences from the available data on electrophonic sounds. This is due primarily to the fact that it is the local transduction conditions that dictate whether or not an electrophonic sound will be heard (Keay, \\cite{KEAY1}, \\cite{KEAY5}). A fireball that some observers report as being electrophonic may be ``silent'' to other near--by witnesses simply because the environmental conditions have changed. Also, personal ``in--field'' experience has revealed that unsuspecting public observers often fail, at least initially, to mention that they heard an associated sound when describing a fireball event, thinking that the sounds were either an illusory or irrelevant coincidences. All this being said, the literature survey conducted by Kaznev (\\cite{KAZNEV}) revealed that from a total of 888 electrophonic meteor events some 76 (8.5~\\%) would qualify for membership in our burster category. The survey by Keay (\\cite{KEAY4}) indicates that 31 (10~\\%) out of the 301 events considered would qualify as bursters. We note also that both of the instrumentally observed electrophonic meteors fall into the short duration burster class. The observations also indicate that electrophonic burster meteors must be very bright. Indeed, the 1981, August 13th electrophonic fireball event recorded in Japan had an estimated visual magnitude of $-6$, while the fireball of the 1993, August 11th had an estimated visual magnitude of $-10$. Keay (\\cite{KEAY1}) and Bronshten (\\cite{BRONSHTEN}) have developed a robust theory to explain the extended ``rushing'' or ``crackling'' electrophonic sounds. The key physical mechanism identified in the production of these sounds is the freezing--in and ``twisting'' of the geomagnetic field in the turbulent wake behind a large meteoroid. In this mechanism it is the release of the strain energy in the geomagnetic field that produces very low frequency (VLF) radio waves and these, depending upon the local environmental conditions, are transducted into audible sounds. The generation of a VLF radio wave signal will proceed provided that the Reynolds number in the meteor ablation column is greater than $10^{6}$ (i.e., the ablation column is turbulent) and that the magnetic Reynolds number is concomitantly greater than unity. Irrespective of the environmental surroundings, if the Reynolds number conditions are satisfied, then an appropriately designed receiver should detect the VLF radio signal. These conditions can be useful in order to evaluate the dimension of the meteoroid (Beech, \\cite{BEECH}). The mechanisms responsible for producing electrophonic bursters have not been as straightforward to annotate as those for the extended sounds. However, the inherent characteristics of burster events suggest that they relate to catastrophic rather than on--going events in the atmospheric ablation of a meteoroid. This phenomenological argument suggests an association between electrophonic bursters and meteor flares and terminal detonations. ", + "conclusions": "In this Letter we have presented a novel and hitherto unreported model for the generation of short--duration electrophonic bursters. The model assumes the catastrophic disruption of a large meteoroid and the subsequent separation of electrons and ions by an energetic shock wave. Since meteoroids enter the Earth's atmosphere with hypersonic velocities, an airburst detonation results in the formation of a shock wave that propagates in the plasma boundary. As a result of the large temperature and pressure gradients across the shock, there is significant diffusion of the electron gas with respect to the ion gas, with the result that an electric field is produced by the space charge separation. It is the rapid variation in the electric field strength that results in the potential generation of electrophonic sounds. At present the observations only afford one case (the August 11th, 1993 fireball event) in which actual measurements can be compared against the predictions. We are pleased to find, however, that in this one case there is an excellent agreement between the predicted electric field strength variation (as described by Eq.~\\ref{e:efield}) and the measurements. While our model was specifically developed to explain the electrophonic burster events, we note that the same basic mechanism may also operate with respect to producing extended electrophonic sounds. That is, if shock waves are produced within the hypersonic flow around a large ablating meteoroid, the space charge separation mechanism can ``run'' in a temporally extended fashion. We hope to investigate this situation in future work." + }, + "0112/astro-ph0112006_arXiv.txt": { + "abstract": "Modified versions of two ``standard'' pulsar search techniques are presented that allow large-scale searches for pulsations in long duration high-energy data sets using relatively modest amounts of computer time. For small numbers of photons ($N_{phot} \\la 10^4$), optimized brute-force epoch folding searches are preferred. For larger numbers of photons, advanced Fourier domain acceleration searches are used. Using these techniques, my collaborators and I have searched \\emph{Chandra} observations of the Cas~A supernova remnant (SNR) point source and the isolated neutron star RX~J1856.5$-$3754 for pulsations, and confirmed the 65.6\\,ms pulsar in the 3C~58 SNR during a blind search of archival \\emph{RXTE} data. ", + "introduction": "Sensitive searches for pulsars in long duration high-energy (i.e. X-ray and $\\gamma$-ray) observations are difficult. For relatively short observations ($T_{obs} \\la 20$\\,ks) all known isolated pulsars can be assumed to have a constant spin period during the observation. In this limiting case, searches for relatively bright pulsars are usually implemented using simple epoch folding or a standard Fourier analysis \\citep[e.g.][]{lde+83,van89}. When $T_{obs}$ is long, however, the spin-down of young high-energy pulsars becomes important. The corresponding change in frequency during an observation can make these pulsars invisible to traditional search techniques. Another difficulty involved in searching very long observations (i.e. many days) is the extraordinarily large number of independent trials required for a blind search. For epoch folding searches that cover a wide range of both frequency $f$ and frequency derivative $\\dot{f}$, the computational complexity is $\\propto N_{phot}T_{obs}^2$, or $\\propto T_{obs}^3$ for a source of constant flux. Advanced Fourier analyses of very long time series (i.e. $(2-20)\\times 10^8$ points where the photons have been binned) have been carried out by a few groups \\citep[e.g.][]{mkk+00, ckl+01} and are, in general, more efficient due to the use of the FFT (computational complexity $\\propto T_{obs}^2\\log T_{obs}$). The difficulties involved in computing many FFTs of this size, however, have relegated such analyses to those with access to large-scale super-computing. The techniques that I describe here allow large-scale blind searches of very long high-energy observations to be carried out on modest (albeit modern) workstations or inexpensive workstation clusters. ", + "conclusions": "" + }, + "0112/astro-ph0112230_arXiv.txt": { + "abstract": "We have obtained 11.7 ${\\mu}$m and 17.9 ${\\mu}$m images at the Keck I telescope of the circumstellar dust emission from L$_{2}$ Pup, one of the nearest ($D$ = 61 pc) mass-losing, pulsating, red giants that has a substantial infrared excess. We propose that the star is losing mass at a rate of ${\\sim}$3 ${\\times}$ 10$^{-7}$ M$_{\\odot}$ yr$^{-1}$. Given its relatively low luminosity (${\\sim}$ 1500 L$_{\\odot}$), relatively high effective temperature (near 3400 K), relatively short period (${\\sim}$ 140 days), and the inferred gas outflow speed of 3.5 km s$^{-1}$, standard models for dust-driven mass loss do not apply. Instead, the wind may be driven by the stellar pulsations with radiation pressure on dust being relatively unimportant, as described in some recent calculations. L$_{2}$ Pup may serve as the prototype of this phase of stellar evolution where it could lose ${\\sim}$15\\% of its initial main sequence mass. ", + "introduction": "Red giants with luminosities near 10$^{4}$ L$_{\\odot}$ typically display infrared excesses that are the result of dust formed in a stellar wind reprocessing optical light emitted by the central star (Habing 1996). The mass loss rates can exceed 10$^{-5}$ M$_{\\odot}$ yr$^{-1}$ and the outflow speed is typically near 15 km s$^{-1}$. These winds are important in stellar evolution because the star may lose over half of its initial main sequence mass during this phase. A standard model to explain this mass loss is that as the stars pulsate, shock waves form in the atmosphere which drive matter above the nominal photosphere. In the post-shock region, the matter cools and grains form. After the dust is created, radiation pressure then expels the matter into the interstellar medium (see, for example, Lamers \\& Cassinelli 1999, Willson 2000). Although the standard model successfully explains the properties of many red giants winds, there are some systems where it does not apply (see Jura \\& Kahane 1999). This ``standard model\" does not consider the effects of binaries, rotation or magnetic fields. In a recent set of calculations of models for mass loss from pulsating red giants, Winters et al. (2000) found a class of solutions, that are distinct from the standard model and which the authors denote as ``case B\", where the mass loss rate is less than 3 ${\\times}$ 10$^{-7}$ M$_{\\odot}$ yr$^{-1}$ at an outflow speed less than 5 km s$^{-1}$. In these models, mechanical energy from the pulsations drives the mass loss; radiation pressure on grains is a relatively minor effect. These ``case B\" models apply to stars of relatively low luminosity, typically where $L_{*}/M_{*}$ $<$ 3500 $L_{\\odot}/M_{\\odot}$, the pulsational periods are less than 300 days and the effective temperatures of the mass-losing stars are greater than 2600 K. Winters et al. (2000) were not able to identify any stars which clearly exhibit ``B-class\" mass loss. Here, we report detailed mid-infrared observations obtained at the 10m Keck telescope of L$_{2}$ Pup, a relatively nearby pulsating, mass-losing red giant. We propose that our data for the circumstellar shell around L$_{2}$ Pup can be best interpreted as a ``B-class\" wind. ", + "conclusions": "We have obtained mid-IR images of the dust around L$_{2}$ Pup. We find the following: \\begin{itemize} \\item{We estimate a total mass loss rate of ${\\sim}$3 ${\\times}$ 10$^{-7}$ M$_{\\odot}$ yr$^{-1}$. During the current phase of its evolution, L$_{2}$ Pup may lose ${\\sim}$15\\% of its initial main sequence mass.} \\item{ Non-radial pulsations may at least in part account for the observed asymmetric mass loss, time variations of the infrared fluxes, and time variations of the position angle of the optical polarization.} \\item{L$_{2}$ Pup might serve as the prototype of an outflow driven by stellar pulsations with radiation pressure on dust being relatively unimportant, as in models described by Winters et al. (2000).} This work has been partly supported by NASA. \\end{itemize}" + }, + "0112/hep-ph0112125_arXiv.txt": { + "abstract": "We show that the temperatures of the emergent non-electron neutrinos and the binding energy released by a galactic Type II supernova are determinable, assuming the Large Mixing Angle (LMA) solution is correct, from observations at the Sudbury Neutrino Observatory (SNO) and at Super-Kamiokande (SK). If the neutrino mass hierarchy is inverted, either a lower or upper bound can be placed on the neutrino mixing angle $\\theta_{13}$, and the hierarchy can be deduced for adiabatic transitions. For the normal hierarchy, neither can $\\theta_{13}$ be constrained nor can the hierarchy be determined. Our conclusions are qualitatively unchanged for the proposed Hyper-Kamiokande detector. \\pacs{} ", + "introduction": " ", + "conclusions": "" + }, + "0112/astro-ph0112556_arXiv.txt": { + "abstract": "The observational determination of the behaviour of the star formation rate (SFR) with look-back time or redshift has two main weaknesses: 1~- the large uncertainty of the dust/extinction corrections, and 2~- that systematic errors may be introduced by the fact that the SFR is estimated using different methods at different redshifts. Most frequently, the luminosity of the \\Ha\\ emission line, that of the forbidden line \\OII\\ and that of the far ultraviolet continuum (UV) are used with low, intermediate and high redshift galaxies respectively. To assess the possible systematic differences among the different SFR estimators and the role of dust, we have compared SFR estimates using \\Ha , SFR(\\Ha), \\OII \\AA , SFR(OII), UV, SFR(UV) and FIR, SFR(FIR) luminosities of a sample comprising the 31 nearby star forming galaxies having high quality photometric data in the UV, optical and FIR. We review the different ``standard\" methods for the estimation of the SFR and find that while the standard method provides good agreement between SFR(\\Ha) and SFR(FIR), both SFR(OII) and SFR(UV) are systematically higher than SFR(FIR), irrespective of the extinction law. We show that the excess in the SFR(OII) and SFR(UV) is mainly due to an overestimate of the extinction resulting from the effect of underlying stellar Balmer absorptions in the measured emission line fluxes. Taking this effect into consideration in the determination of the extinction brings the SFR(OII) and SFR(UV) in line with the SFR(FIR) and simultaneously reduces the internal scatter of the SFR estimations. Based on these results we have derived ``unbiased\" SFR expressions for the SFR(UV), SFR(OII) and SFR(\\Ha). We have used these estimators to recompute the SFR history of the Universe using the results of published surveys. The main results are that the use of the unbiased SFR estimators brings into agreement the results of all surveys. Particularly important is the agreement achieved for the SFR derived from the FIR/mm and optical/UV surveys. The ``unbiased\" star formation history of the Universe shows a steep rise in the SFR from $z=0$ to $z=1$ with SFR $\\propto (1+z)^{4.5}$ followed by a decline for $z>2$ where SFR $\\propto (1+z)^{-1.5}$. Galaxy formation models tend to have a much flatter slope from $z=0$ to $z=1$. ", + "introduction": "The knowledge of the history of star formation at cosmic scales is fundamental to the understanding of the formation and evolution of galaxies. Madau and collaborators combined the results from the ultraviolet surveys of Lilly \\etal \\shortcite{1996Lilly} with the information from the Hubble Deep Field to give an estimate of the star formation history from $z = 0$ to $z = 4$. Subsequent studies have indicated the crucial role played by dust in the estimates of SFR. Large correction factors were suggested for $z > 1$--$2$ by several authors (Meurer \\etal 1997, Meurer, Heckman and Calzetti 1999, Steidel \\etal 1999, Dickinson 1998). But even these large dust extinction corrections do not seem to be enough to bring the optical/UV SFR estimates in line with the mm/sub-mm ones (Hughes \\etal 1998; Rowan-Robinson \\etal 1997 (RR97); Chapman \\etal 2001). On top of the uncertainties associated with the extinction correction, most of the SFR estimates have been performed using expressions derived from spectra constructed using population synthesis methods, an approach that requires four rather uncertain ingredients: 1) an initial mass function (IMF); 2) a stellar evolutionary model grid giving the luminosity and effective temperature as a function of time; 3) a stellar atmospheres grid that assigns a spectrum to each star for a given luminosity and effective temperature, and 4) a star formation history. The fact that the redshift evolution of the SFR is constructed using different estimators at different redshift ranges is a potential source of systematic effects with redshift that can distort the shape of the evolutionary curve. In practice the \\Ha\\ luminosity is used to estimate the SFR for galaxies with redshifts up to 0.4; the [OII]~$\\lambda$3727\\AA\\ line , for those with 0.4$>$z$>$1.0 and the UV continuum luminosity for galaxies with z$>$2.0. In addition, dust extinction corrections are not treated uniformly over the whole redshift range. It is therefore important to ensure that there are no systematic differences between the different estimators and corrections that can distort the results. This paper has two main aspects, in the first 5 sections we review the ``standard\" methods for the estimation of the SFR and test the consistency of the different SFR estimators by applying them to a sample of well studied nearby star forming galaxies and comparing the results. In the absence of systematic differences among them, all should give the same SFR for each one of the galaxies in the sample. We then use the results of the nearby sample to construct a set of ``unbiased'' SFR estimators and apply them to published surveys. ", + "conclusions": "One aspect that has generated many discussions regarding the SF history diagram in Cosmology, is the lack of confidence on the reddening corrections. This is compounded with the low level of agreement found until now between the optical/UV and FIR determinations of the SFR. In the first part of this paper, we have investigated the possible systematic differences between SFR estimators by applying them to a sample of nearby star forming galaxies with good photometric data from the UV to the FIR. We found that the main source of systematic differences among the SFR rate estimators is related to the presence of stellar Balmer absorption in the spectrum of emission line galaxies. The main effect of the Balmer absorptions is to produce an overestimate of the reddennnig when their effect is not included. We showed that taking into account the underlying Balmer absorptions effect in the estimates of reddening, removes most of the systematic differences between the SFR estimators in the optical/UV and FIR. Furthermore the scatter of the SFR estimations is considerably reduced by the application of the corrections. We also found that the escape of photons plays a minor role compared to that of the Balmer absorptions. These results give renewed confidence to the estimates of SFR for star forming galaxies in general and for samples similar to the one presented here in particular. {\\bf Thus, our central result is that the extinction correction including the effects of an underlying stellar Balmer absorption brings into agreement all four SFR estimators, and that the photon escape correction seems to play a minor role.} In the second part of the paper we used the average results for our sample to construct a set of ``unbiased'' SFR estimators. These ``unbiased'' SFR estimators expressions include statistically the underlying Balmer absorption and photon escape corrections to the extinction estimates and bring the four SFR estimators studied here into the same system. We thus obtained consistent results between the SFR estimators in the optical/UV and FIR. The application of these ``unbiased'' SFR estimators to a compilation of surveys has produced a SFR history of the universe where all surveys results agree whitin the errors. Particularly important is the level of agreement achieved between the FIR/mm and optical/UV SFR results. Our ``new\" and unbiased SFR history of the universe shows a steep rise in the SFR rate from $z=0$ to $z=1$ with SFR $\\propto (1+z)^{4.5}$ followed by a mild decline for $z>2$ where SFR $\\propto (1+z)^{-1.5}$. The steep increase to $z=1$ seems in line with recent determinations of the SFR using only the \\Ha\\ estimator. Most galaxy formation models tend to have a much flatter slope from $z=0$ to $z=1$." + }, + "0112/astro-ph0112283_arXiv.txt": { + "abstract": "We describe the approach adopted by our group to derive the star formation history (SFH) of the chosen LMC field, as part of the experiment to compare the predictions obtained by different groups with the synthetic Colour-Magnitude diagram method. We point out what are the evolutionary characteristics of the observed stellar populations, and present the SFH scenario which appears to better account for them. We emphasize the importance of adopting a {\\it critical} approach when dealing with this method. ", + "introduction": "The use of synthetic Colour-Magnitude diagrams (CMDs) allows people to derive the SF history of nearby galaxies by interpreting the observed features of their resolved stellar populations in terms of stellar evolution (Tosi et al 1991). Nowadays, this approach is widely followed by the international community, but some people still wonder whether or not the method and stellar evolution theories are reliable enough to guarantee consistency between results obtained by different groups and/or procedures. The Coimbra experiment was set up to compare the results obtained by different groups from the same set of high quality photometric data, treated with the same procedures both for the data reduction and for the artificial star tests. A general report on the experiment is given by Skillman \\& Gallart (this volume, hereafter SG). The data set refers to HST-WFPC2 images of a field in the LMC bar, kindly made available by T. Smecker-Hane. Our group has adopted the data catalogue provided by A. Dolphin from data reduction performed with the HSTphot package, and his catalogue for the results of the artificial star tests. ", + "conclusions": "" + }, + "0112/astro-ph0112410_arXiv.txt": { + "abstract": "We have found a bar of shocked molecular hydrogen (\\h) towards the \\oh maser located at the projected intersection of supernova remnant (SNR) G359.1--0.5 and the nonthermal radio filament, known as the Snake. The \\h\\ bar is well aligned with the SNR shell and almost perpendicular to the Snake. The \\oh maser is located inside the sharp western edge of the \\h\\ emission, which is consistent with the scenario in which the SNR drives a shock into a molecular cloud at that location. The spectral--line profiles of \\co, \\hcop\\ and CS towards the maser show broad--line absorption, which is absent in the \\tco\\ spectra and most probably originates from the pre--shock gas. A density gradient is present across the region and is consistent with the passage of the SNR shock while the \\h\\ filament is located at the boundary between the pre--shocked and post-shock regions. ", + "introduction": "\\oh maser emission detected towards supernova remnants (SNRs) has been attributed to shock waves driven into adjacent molecular clouds \\citep{frail94}. The production of this maser emission, without the simultaneous production of maser emission in the other three ground state OH transitions at 1612, 1665 and 1667~MHz, requires specific conditions (gas density \\nh$\\sim$\\e{5}\\cm{-3}, gas temperature \\tkin$\\sim$50 -- 125~K, OH column density $N_{\\rm OH}\\sim$ \\e{15} -- \\e{16}\\cm{-2}) which can be attained in the cooling gas behind a non-dissociative shock wave \\citep*{lock99,wardle99}. Surveys by \\citet{frail96}, \\citet{green97} and \\citet{kor98} have found that about 10 per cent of $\\sim$180 observed Galactic SNRs are associated with \\oh masers. The detection fraction increases to 30 per cent for objects in the Galactic Centre region \\citep{yz96-1}. Perhaps the most striking result of the survey towards the Galactic Centre region is the detection of OH masers along the radio--continuum shell of the SNR G359.1--0.5 \\citep{yz95,yz96-2}. The brightest of the masers (maser {\\em A} from \\citealt{yz95}) occurs where a nonthermal filament known as `the Snake' appears to cross the western edge of the SNR shell. The filament differs from other Galactic Centre nonthermal filaments in showing a number of kinks along its 20 arcmin extent \\citep{grey95}, and its origin is not well established \\citep{nicholls95,benford97,uchida96}. The latest model by \\citet{bi00} proposes that the Snake is a magnetic flux tube anchored in dense rotating material. It has been suggested that the Snake is interacting with the shell of G359.1--0.5 because both objects show change in brightness at the apparent crossing point \\citep{uchida92-2,grey95}. This could, however, be an artifact due to superposition of disparate components along the line of sight. Nevertheless, the presence of an \\oh maser -- a signature of SNR/molecular cloud interaction -- at this location is quite intriguing. In order to investigate the proposed interaction between the Snake and the SNR, and to characterise the ambient molecular gas in the region, we observed molecular hydrogen (\\h) and a number of other molecular species, particularly searching for signatures of shocked gas. The observations are discussed in section 2, the results are presented in section 3 and discussed in section 4. The conclusions are presented in section 5. ", + "conclusions": "We carried out near--IR and millimetre wavelength observations of the region around the \\oh maser located at the apparent intersection of the nonthermal filament the Snake and the SNR shell G359.1--0.5. A bar of \\h\\ emission was found encompassing the \\oh maser and aligned with the SNR shell. We suggest that \\h\\ emission originates from the expansion of the SNR blast wave, which is evident from the sharp western edge that corresponds to the forward shock. This shocked \\h\\ emission, as inferred from the 1--0 and 2--1 S(1) line ratio, supports the notion that the \\oh masers association with SNRs are produced in molecular shock waves. Emission from the molecular species \\co, CS, \\hcop\\ and \\tco\\ was detected at the maser velocity. The spectra of the former three species are affected by broad--line absorption which originates from a colder layer of molecular gas. Optically thin \\tco\\ was unaffected by the absorption and is produced mostly in this colder layer, identified as pre--shocked gas. The inferred density gradient across the region places higher density post--shocked gas west of the \\h\\ bar. The distribution of the pre--shock and post--shock gas is consistent with the passage of the SNR shock and the \\h\\ bar being the boundary between the two regions. The warm molecular gas from post--shocked gas was not detected in millimetre wavelengths probably because of beam dilution. For a better morphological and qualitative study of the shocked molecular gas towards maser {\\em A} we plan to obtain observations at sub--millimetre wavelengths, which will have improved spatial resolution. Observations at other maser positions in the SNR are in progress to obtain a global view of the interaction of the SNR G359.1--0.5 with its molecular gas environment." + }, + "0112/astro-ph0112318_arXiv.txt": { + "abstract": "The \\asca calibration has evolved considerably since launch and indeed, is still evolving. There have been concerns in the literature that changes in the \\asca calibration have resulted in the \\fekalfa lines in active galaxies (AGN) now being {\\it systematically} narrower than was originally thought. If this were true, a large body of \\asca results would be impacted. In particular, it has been claimed that the broad red wing (when present) of the \\fekalfa line has been considerably weakened by changes in the \\asca calibration. We demonstrate explicitly that changes in the \\asca calibration over a period of about eight years have a negligible effect on the width, strength, or shape of the \\fekalfa lines. The reduction in both width and equivalent width is only $\\sim 8\\%$ or less. We confirm this with simulations and individual sources, as well as sample average profiles. The average profile for type 1 AGN is still very broad, with the red wing extending down to $\\sim 4$ keV. The reason for the claimed, apparently large, discrepancies is that in some sources the \\fekalfa line is complex, and a single-Gaussian model, being an inadequate description of the line profile, picks up different portions of the profile with different calibration. Single-Gaussian fits do not therefore model all of the line emission in some sources, in which case they do not compare old and current calibration since the models do not then describe the data. ", + "introduction": "The {\\it Advanced Satellite for Cosmology and Astrophysics} (\\asca -- Tanaka, Inoue, \\& Holt 1994) found that the \\fekalfa fluorescent emission line in Seyfert~1 galaxies is often very broad, and this is generally interpreted as the result of an origin in matter in an accretion disk rotating around a central black hole (see Fabian \\etal 2000 and references therein). The line profile is thought to be sculpted by characteristic gravitational and Doppler energy shifts. Currently, study of the \\fekalfa emission line is the only way to probe matter within a few to tens of gravitational radii of a black hole. Indeed, the broad \\fekalfa lines provide some of the strongest evidence to date for the existence of black holes. It has been known since the early days of \\asca that the \\fekalfa line in AGN is not always tremendously broad (e.g. Ptak 1994) and that the profile comes in a variety of shapes (e.g. Yaqoob \\& Weaver 1996; Nandra \\etal 1997, hereafter N97). It has also been known that Seyfert type 1.9--2.0 galaxies have a predominantly narrow \\fekalfa line, with FWHM less than $\\sim 10,000$ km/s, probably originating in cold matter far from the black hole (e.g. NGC 2992; Weaver \\etal 1996). It was apparent that many Seyfert 1.5--1.9 galaxies likely have composite narrow and broad \\fekalfa line components (e.g. Weaver \\etal 1993, Yaqoob \\etal 1995, Weaver \\etal 1997, Weaver \\& Reynolds, 1998). \\chandra and \\xmm have now shown that composite narrow and broad \\fekalfa lines are common even in Seyfert 1 galaxies and quasars (Yaqoob \\etal 2001a, Kaspi \\etal 2001, Reeves \\etal 2000, Pounds \\etal 2001). Recently Lubi\\'{n}ski and Zdziarski (2001; hereafter LZ01) have claimed that the \\fekalfa lines in AGN are {\\it systematically} narrower and weaker (smaller equivalent width, or EW) than had previously been thought. They speculated that this result can be attributed to changes in the calibration of the \\asca instruments. However they did not attempt to demonstrate this by isolating calibration effects. The purpose of the present paper is to explicitly demonstrate that changes in the \\asca calibration have {\\it not} resulted in any significant changes in the width, equivalent width, or shape of the \\fekalfa lines in AGN. This issue is important to resolve because speculation about the \\asca calibration has brought into question a significant body of \\asca results and has rendered the astronomical community rather confused. We will also give the real explanation of the results of LZ01, which turns out to be a very subtle data analysis issue. ", + "conclusions": "\\label{concl} We find that the width, equivalent width, and shape of the \\fekalfa lines in AGN have not been significantly affected by changes in the \\asca calibration over a period of more than eight years. In particular, the width and equivalent width change by only $\\sim 8\\%$ or less. The red wing, in the case of a highly skew \\fekalfa line profile as in \\mcg is reduced by a similar magnitude. In most cases the statistical errors are larger, so the changes in \\asca calibration are inconsequential for almost all \\asca observations. We also compared the sample composite \\fekalfa line profiles using OLD and CURRENT calibration and find that the differences are still not statistically significant. In some individual cases the \\fekalfa lines are complex so single-Gaussian models are inadequate for describing the profile and consequently only model a portion of the line. Then, the line width and equivalent width may be seriously understimated, leading to erroneous conclusions about the astrophysics of the line and about instrument calibration. When adequately modeled, we again find no significant difference in the \\fekalfa line parameters with respect to changes in calibration. The authors acknowledge support from NASA grants NCC-5447 (TY, UP), NAG5-10769 (TY), and NAG5-7067 (KN). This research made use of the HEASARC online data archive services, supported by NASA/GSFC. It also made use of the {\\sc tartarus} AGN database which is supported by Jane Turner and Kirpal Nandra under NASA grants NAG5-7385 and NAG5-7067. Special thanks to Prof. Y. Tanaka for his inputs and motivation. The authors also wish to thank Koji Mukai, Ken Ebisawa, Keith Arnaud, and other members of the ASCATEAM for their input to this work. They also thank Jane Turner, Ian George, Kim Weaver for helpful discussions and comments. We are grateful to Andrzej Zdziarski, the referee, and Piotr Lubi\\'{n}ski, for making suggestions to greatly improve this paper. \\newpage \\begin{deluxetable}{lcccccccc} \\tablecaption{Single and Double-Gaussian Models for NGC 3783, Mkn 841, and NGC 3516} \\tablecolumns{9} \\tablewidth{0pt} \\tablehead{ \\colhead{} & \\colhead{$E_{N}$} & \\colhead{$\\sigma_{N}$} & \\colhead{$EW_{N}$} & \\colhead{$E_{B}$} & \\colhead{$\\sigma_{B}$} & \\colhead{$EW_{B}$} & \\colhead {$\\chi^{2}$(d.o.f)} & \\colhead{$\\Delta \\chi^{2 \\ \\dagger}$} \\nl \\colhead{} & \\colhead{(keV)} & \\colhead{(keV)} & \\colhead{(eV)} & \\colhead{(keV)} & \\colhead{(keV)} & \\colhead{(eV)} & \\colhead{ } & \\colhead{} \\nl } \\startdata NGC3783(1) & & & & & & & & \\nl S (OLD) & $^{}_{}$ & $^{}_{}$ & $^{}_{}$ & $6.13^{+0.16}_{-0 .19}$ & $0.69^{+0.36}_{-0.27}$ & $527^{+346}_{-219}$ & $804.5(7 84)$ & $ $ \\nl S (CURRENT) & $6.41^{+0.16}_{-0.50}$ & $0.0^{+1.09}_{-0.00}$ & $169^{+45}_{-40}$ &$^{}_{}$ & $^{}_{}$ & $^{}_{}$ & $806.7(7 84)$ & $ $ \\nl D (OLD) & $6.41^{+0.05}_{-0.06}$ & $0.001^{}_{}(f)$ & $114^{+44}_{-59}$ & $5.88^{+0.32}_{-0.41}$ & $0.91^{+0.91}_{-0 .44}$ & $469^{+761}_{-253}$ & $777.5(782)$ &$27.0$ \\nl D (CURRENT) & $6.41^{+0.54}_{-0.55}$ & $0.001^{}_{}(f)$ & $129^{+54}_{-54}$ & $5.88^{+0.36}_{-0.44}$ & $1.35^{+0.71}_{-0 .30}$ & $700^{+474}_{-314}$ & $777.3(782)$ & $29.4$ \\nl Mkn 841 & & & & & & & & \\nl S (OLD) & $^{}_{}$ & $^{}_{}$ & $^{}_{}$ & $6.11^{+0.41}_{-0 .26}$ & $0.70^{+0.35}_{-0.66}$ & $738^{+575}_{-355}$ & $422.6(3 98)$ & $ $ \\nl S (CURRENT) & $6.48^{+0.09}_{-0.08}$ & $0.09^{+0.79}_{-0.09}$ & $257^{+116}_{-101}$ & $^{}_{}$ & $^{}_{}$ & $^{}_{}$ & $419.8( 398)$ & $ $ \\nl D (OLD) & $6.51^{+0.10}_{-0.12}$ & $0.001^{}_{}(f)$ & $154^{+128}_{-102}$ & $5.87^{+0.37}_{-0.49}$ & $0.72^{+0.56}_{ -0.48}$ & $556^{+672}_{-288}$ & $407.2(396)$ & $15.4$ \\nl D (CURRENT) & $6.51^{+0.11}_{-0.09}$ & $0.001^{}_{}(f)$ & $250^{+96}_{-178}$ & $5.64^{+0.67}_{-0.42}$ & $0.59^{+1.09}_{- 0.30}$ & $324^{+885}_{-231}$ & $406.1(396)$ & $13.7$ \\nl NGC3516 & & & & & & & & \\nl S (OLD) & $^{}_{}$ & $^{}_{}$ & $^{}_{}$ & $6.09^{+0.12}_{-0 .14}$ & $0.65^{+0.22}_{-0.19}$ & $360^{+143}_{-106}$ & $873.4(7 70)$ & $ $ \\nl S (CURRENT) & $6.33^{+0.07}_{-0.22}$ & $0.23^{+0.60}_{-0.10}$ & $171^{+245}_{-45}$ & $^{}_{}$ & $^{}_{}$ & $^{}_{}$ & $884.3(7 70)$ & $ $ \\nl D (OLD) & $6.41^{+0.04}_{-0.05}$ & $0.001^{}_{}(f)$ & $75^{+22}_{-40}$ & $5.86^{+0.36}_{-0.22}$ & $0.77^{+0.38}_{-0 .11}$ & $294^{+233}_{-84}$ & $840.6(768)$ & $32.8$ \\nl D (CURRENT) & $6.41^{+0.08}_{-0.04}$ & $0.001^{}_{}(f)$ & $79^{+30}_{-34}$ & $5.91^{+0.37}_{-0.33}$ & $1.00^{+0.93}_{-0 .24}$ & $318^{+581}_{-115}$ & $847.8(768)$ & $36.5$ \\nl \\enddata \\vspace{-5mm} \\tablecomments{ Comparison of the OLD and CURRENT calibration for single (S) and double-Gaussian (D) models of the \\fekalfa line. Subscripts $N$ and $B$ denote the parameters (center energy, $E$, intrinsic width, $\\sigma$, and equivalent width, $EW$) for the narrow and broad \\fekalfa line components respectively. Statistical errors are 68\\% confidence for four ($\\Delta \\chi^{2} = 4.70$) and six ($\\Delta \\chi^{2} = 7.00$) interesting parameters, for the single and double-Gaussian models respectively. $^{(\\dagger)}$ $\\Delta \\chi^{2}$ is the difference in $\\chi^{2}$ between single and double-Gaussian fits. } \\end{deluxetable}" + }, + "0112/astro-ph0112404_arXiv.txt": { + "abstract": "NGC4826 (M64) is a nearby Sab galaxy with an outstanding, absorbing dust lane (called the ``Evil Eye'') asymmetrically placed across its prominent bulge. In addition, its central region is associated with several regions of ongoing star formation activity. We obtained accurate low-resolution ($\\rm 4.3~\\AA$/pixel) long-slit spectroscopy (KPNO 4-Meter) of NGC4826 in the 5300 -- $\\rm 9100~\\AA$ spectral range, with a slit of 4.4 arcmin length, encompassing the galaxy's bulge size, positioned across its nucleus. The wavelength-dependent effects of absorption and scattering by the dust in the Evil Eye are evident when comparing the observed stellar spectral energy distributions (SEDs) of pairs of positions symmetrically located with respect to the nucleus, one on the dust lane side and one on the symmetrically opposite side of the bulge, under the assumption that the intrinsic (i.e. unobscured) radiation field is to first order axi-symmetric. We analyzed the SED ratios for a given number of pairs of positions through the multiple-scattering radiative transfer model of Witt \\& Gordon. As a main result, we discovered strong residual Extended Red Emission (ERE) from a region of the Evil Eye within a projected distance of about 13 arcsec from the nucleus, adjacent to a broad, bright HII region, intercepted by the spectrograph slit. ERE is an established phenomenon well-covered in the literature and interpreted as originating from photoluminescence by nanometer-sized clusters, illuminated by UV/optical photons of the local radiation field. In the innermost part of the Evil Eye, the ERE band extends from about 5700 to $\\rm 9100~\\AA$, with an estimated peak intensity of $\\rm \\sim 3.7 \\times 10^{-6}~ergs~s^{-1}~\\AA^{-1}~cm^{-2}~sr^{-1}$ near $\\rm 8300~\\AA$ and with an ERE-to-scattered light band-integrated intensity ratio, $I(ERE)/I(sca)$, of about 0.7. At farther distances, approaching the broad, bright HII region, the ERE band and peak intensity shift toward longer wavelengths, while the ERE band-integrated intensity, $I(ERE)$, diminishes and, eventually, vanishes at the inner edge of this HII region. The radial variation of $I(ERE)$ and $I(ERE)/I(sca)$ does not match that of the optical depth of the model derived for the dust lane. By contrast, the radial variation of $I(ERE)$, $I(ERE)/I(sca)$ and of the ERE spectral domain seems to depend strongly on the strength and hardness of the illuminating radiation field. In fact, $I(ERE)$ and $I(ERE)/I(sca)$ diminish and the ERE band shifts toward longer wavelengths when both the total integrated Lyman continuum photon rate, $Q(H^0)_{TOT}$, and the characteristic effective temperature, $T_{eff}$, of the illuminating OB-stars increase. $Q(H^0)_{TOT}$ and $T_{eff}$ are estimated from the extinction-corrected $\\rm H \\alpha~(\\lambda = 6563~\\AA)$ line intensity and line intensity ratios $[NII] (\\lambda 6583)/H \\alpha$ and $[SII] (\\lambda \\lambda 6716+6731)/H \\alpha$, respectively, and are consistent with model and observed values typical of OB-associations. Unfortunately, we do not have data shortward of $\\rm 5300~\\AA$, so that the census of the UV/optical flux is incomplete. The complex radial variation of the ERE peak intensity and peak wavelength, of $I(ERE)$ and $I(ERE)/I(sca)$ with optical depth and strength of the UV/optical radiation field is reproduced in a consistent way through the theoretical interpretation of the photophysics of the ERE carrier by Smith \\& Witt, which attributes a key-role to the experimentally established recognition that photoionization quenches the luminescence of nanoparticles. When examined within the context of ERE observations in the diffuse ISM of our Galaxy and in a variety of other dusty environments such as reflection nebulae, planetary nebulae and the Orion Nebula, we conclude that the ERE photon conversion efficiency in NGC4826 is as high as found elsewhere, but that the size of the actively luminescing nanoparticles in NGC4826 is about twice as large as those thought to exist in the diffuse ISM of our Galaxy. ", + "introduction": "The ``Evil Eye'' Galaxy (NGC4826; M64) is a relatively isolated, nearby Sab galaxy, located at RA(B1950.0) $\\rm 12h~54m~16.4s$ and Dec(B1950.0) $\\rm +21d~57m~05s$, according to de Vaucouleurs et al. (1991 - RC3). Its distance from us is uncertain by a factor of 2: e.g. Rubin (1994) uses a rounded 5 Mpc distance, as a compromise between the quoted low and high distances (Braun, Walterbos \\& Kennicutt 1992; Tully 1988; Kraan-Korteweg \\& Tammann 1980). Hereafter we adopt the distance quoted by Rubin. NGC4826 has an optical size of $\\rm 10.0 \\times 5.4~arcmin^2$ (RC3) and position angle of the line of nodes (PA for short) of $\\rm 115^o$, on the basis of the outer contours on the POSS (Nilson 1973; see also Rubin 1994). According to the latter author, the red isophotal diameters imply an inclination of $\\rm 60^o$. We refer the reader to Rubin (1994) for further details about morphology and geometry of the Evil Eye Galaxy. NGC4826 owes its surname to the presence of an outstanding, absorbing dust lane (the Evil Eye), asymmetrically placed across its prominent bulge. The central region of NGC4826 shows a further peculiarity, since it is associated with several HII regions (J. Young 2000, private communication; see also Keel 1983a; Rubin 1994). At the Evil Eye, NICMOS imaging in the $\\rm Pa \\alpha$ ($\\rm \\lambda = 1.8751~\\mu m$) line, obtained by B\\\"oker et al. (1999), reveals the presence of an inner arc-shaped region of ongoing star formation activity and of a broad lane of widely distributed, bright HII regions as well, in agreement with the morphology of the $\\rm H \\alpha$ ($\\rm \\lambda = 6563~\\AA$) line emission (J. Young 2000, private communication; see also Keel 1983a; Rubin 1994). Spectroscopic studies (e.g. Keel 1983b; Alonso-Herrero et al. 2000) classify this galaxy as a starburst-powered LINER. Sakamoto et al. (1999) imaged the central region of NGC4826 in the $\\rm ^{12} CO (J = 1 - 0)~(\\lambda = 2.6~mm)$ line at the Nobeyama Millimeter Array, with a $\\rm 4.6 \\times 4.3~arcsec^2$ synthesized beam and an $\\rm 8.1~km~s^{-1}$ channel width. They find that the average hydrogen column density of the molecular clouds at the Evil Eye is about $\\rm 1-3 \\times 10^2~M_{\\odot}~pc^{-2}$, 2 times larger than on the opposite side of the nucleus, by adopting the conventional CO-to-$\\rm H_2$ conversion factor $\\rm X = 3.0 \\times 10^{20}~cm^{-2}~(K~km~s^{-1})^{-1}$, determined in the Galactic disk (Scoville et al. 1987; Solomon et al. 1987; Strong et al. 1988). This average hydrogen column density produces a V-band ($\\rm \\lambda = 0.55~\\mu m$) extinction of $\\rm A_V = 7$ -- 21 mag, if the dust is uniformly distributed and a Galactic dust-to-gas ratio is assumed, but the net asymmetry in extinction, $\\rm \\Delta A_V$, due to the excess gas in the Evil Eye, is only of about 3 -- 10 mag (Sakamoto et al. 1999). As noted by the latter authors, significant deviations of the CO-to-$\\rm H_2$ conversion factor from the standard value (e.g. Wilson \\& Scoville 1990) do exist within our Galaxy and in different starburst galaxies, nuclei of galaxies and low-metallicity galaxies (e.g. Maloney \\& Black 1988; Nakai \\& Kuno 1995; Israel 1997). A value of $\\rm X = 1.06~(\\pm 0.14) \\times 10^{20}~cm^{-2}~(K~km~s^{-1})^{-1}$, derived by Digel, Hunter \\& Mukherjee (1995) in Orion, sounds more reasonable to us in the case of the dust lane of NGC4826, given the presence of several HII regions there. In this case, the CO measurement of Sakamoto et al. (1999) gives $\\rm \\Delta A_V = 1$ -- 3.5 mag. Sakamoto et al. also find that the velocity field is quite ordinary with a hint of noncircular motions at the radius of four CO clumps, three of the clumps being associated with the dust lane. By contrast, two nested counterrotating HI gas disks, at 1 kpc and in the outer region of the stellar disk, have been discovered in this galaxy (Braun, Walterbos \\& Kennicutt 1992; Rix et al. 1995). As quoted by Sakamoto et al., the quiescent velocity field suggests that the Evil Eye does not have much effect on the gas kinematics in the central kiloparsec (i.e. the central 50 arcsec, assuming the distance of 4.1 Mpc adopted by these authors). CO line profiles are single-peaked in the Evil Eye, suggesting to Sakamoto et al. that most of the molecular gas in the dark lane has already settled in the disk even if the gas accreted from outside. This conclusion is at variance with that of Block et al. (1994). The latter authors imaged NGC4826 in the V-band and $\\rm K^{\\prime}$-band ($\\rm \\lambda = 2.1~\\mu m$). The presence of the Evil Eye produces a striking asymmetry in the V-band surface brightness profile, the axial symmetry being restored at $\\rm 2.1~\\mu m$. From the analysis of these surface brightness profiles, oriented through the galaxy's center but sequentially spaced 10 degrees each in position angle, Block et al. determined a least-squares fit value of $\\Delta V / E(V - K^{\\prime}) = 1.04 \\pm 0.05$ for the Evil Eye, by assuming axial symmetry for both the unobscured V-band surface brightness profile and the unreddened $\\rm V - K^{\\prime}$ color profile. The multiple-scattering radiative transfer models of Witt, Thronson \\& Capuano (1992) allowed Block et al. to identify this value with the geometry of a foreground dusty screen with an average V-band optical depth $\\tau_V \\simeq 2.1$. The latter value is equivalent to $\\rm \\Delta A_V = 2.3$ mag, consistent with our previous estimate. Finally, in the analysis of Block et al., the screen model was also supported by the value of the far-IR-to-optical luminosity ratio of NGC4826 and by the agreement between computed and measured (molecular$+$atomic) hydrogen column density. On the basis of the previous arguments, it is reasonable to assume that a moderate amount of gas and dust distributed in a foreground screen is responsible both for the observed amount of attenuation and for the absence of a remarkable kinematic signature in the velocity field at the dust lane of NGC4826. In our opinion, this scenario is supported by the relative weakness of the average CO surface density in the Evil Eye (Sakamoto et al. 1999), compared to the high optical opacity there (Block et al. 1994). In fact, a screen geometry associated with a homogeneous dust distribution (like in models of Witt, Thronson \\& Capuano 1992) gives a lower limit to the computed mass of the dust responsible for a given (observed) reddening (e.g. Witt, Thronson \\& Capuano 1992). The clumpy structure of the dust is manifested by the presence of dark filaments in the NICMOS image of NGC4826 in the $\\rm Pa \\alpha$ line (B\\\"oker et al. 1999). This confirms previous suggestions on the nature of the gaps present in the distribution of the $\\rm H \\alpha$ emission from several HII regions at the dust lane (J. Young 2000, private communication; Keel 1983a; Rubin 1994). This wealth of information makes NGC4826 a suitable test for radiative transfer models. The strong effect of a homogeneous vs. clumpy dust distribution in different star/dust geometries has been highlighted and carefully investigated recently by the new multiple-scattering radiative transfer models of Witt \\& Gordon (2000 -- hereafter referred to as WG00). WG00 have computed the fractions of escaping/scattered radiation in the 1000 -- 30000 $\\rm \\AA$ spectral range from three representative types of galactic environments, filled with either homogeneous or two-phase clumpy dust distributions. Furthermore, they considered two types of interstellar dust, similar to those found in the average diffuse interstellar medium (ISM) of the Milky Way (MW) and in the Bar of the Small Magellanic Cloud (SMC). For a given optical depth of the model, absolute estimates of the escaping radiation require the knowledge of the spectral energy distribution (SED) of the intrinsic (i.e. non attenuated) interstellar radiation field (ISRF), e.g. through the comparison between data and calculations of stellar population synthesis models. In the case of NGC4826, low-resolution spectroscopy of the central region, with a slit encompassing equally extended regions on the obscured and unobscured side of the nucleus, provides to first order an observational determination of the intrinsic stellar SED in the dust lane, by assuming axial symmetry for the ISRF before dust absorption and scattering (cf. Block et al. 1994). As a consequence, it also provides a measurement of the attenuation of the intrinsic ISRF in the Evil Eye. A least-squares fit of the computed fractions of transmitted and scattered radiation at several different wavelengths to the ratio of the observed and intrinsic SEDs, for a given position across the dust lane, identifies the corresponding optical depth. As a further goal, we want to search for a residual dust emission feature, which manifests itself in the familiar broad 0.7 $\\rm \\mu m$ band, known as Extended Red Emission (ERE). Broad-band ($\\rm \\Delta \\lambda \\sim 0.1~\\mu m$) ERE, with a peak intensity wavelength $\\lambda_p(ERE)$ between 0.65 and 0.88 $\\mu m$, has been reported in many different dusty astrophysical environments of our Galaxy, e.g. the diffuse ISM (Gordon, Witt, \\& Friedmann 1998), cirrus clouds (Szomoru \\& Guhathakurta 1998), reflection nebulae (Witt \\& Boroson 1990), planetary nebulae (Furton \\& Witt 1990, 1992), the Orion Nebula (Perrin \\& Sivan 1992) and the high-$b$ dark nebula Lynds 1780 (Chlewicki \\& Laureijs 1987; Mattila 1979). We refer the reader to e.g. Witt, Gordon \\& Furton (1998) for a review of the observed properties of ERE. Here, it is sufficient to say that the ERE is commonly interpreted as a photoluminescence phenomenon originating in dusty environments illuminated by UV/optical photons. Identifying ERE in NGC4826 would extend our current knowledge of ERE to galaxies different from our own, so far limited to the halo of the starburst galaxy M82 (Perrin, Darbon, \\& Sivan 1995) and the 30-Doradus HII region in the Large Magellanic Cloud (LMC -- Darbon, Perrin \\& Sivan 1998), but also provide the theory of the ERE carrier (Witt, Gordon \\& Furton 1998 and references therein) and photophysics (Smith \\& Witt 2001) with a further bench-mark. We present new optical spectroscopy of NGC4826 in Sect. 2. Data reduction and analysis are discussed in Sect. 3. The results of our study are illustrated in Sect. 4. In Sect. 5 we discuss these results within the context of ERE observations in different dusty environments of our Galaxy, laboratory studies of photoluminescence properties of silicon nanocrystals and the models of Smith \\& Witt (2001). Conclusions are summarized in Sect. 6. ", + "conclusions": "NGC4826 is a nearby Sab galaxy with an outstanding asymmetrically placed, absorbing dust lane (called the ``Evil Eye'') across its prominent bulge, associated with several regions of ongoing star formation. For this galaxy, we obtained accurate low-resolution ($\\rm 4.3~\\AA$/pixel) long-slit spectroscopy (KPNO 4-Meter) in the 5300 -- $\\rm 9100~\\AA$ wavelength range, with a slit of 4.4 arcmin length, encompassing the galaxy's bulge size, positioned across its nucleus. This allowed us to study the wavelength-dependent effects of absorption and scattering by the dust by comparing the stellar spectral energy distributions of pairs of positions across the bulge, symmetrically placed with respect to the nucleus, under the assumption that the intrinsic (i.e. unobscured) radiation field is axi-symmetric, except for the extra contribution due to ongoing star formation activity associated with the dust lane. In this analysis we made use of the multiple-scattering radiative transfer calculations of Witt \\& Gordon (2000). As a main result, strong residual Extended Red Emission (ERE) was detected from a region of the Evil Eye within about 13 arcsec projected distance from the NGC4826 nucleus, adjacent to a broad, bright HII region, intercepted by the spectrograph slit. In the innermost portion of the dust lane, the ERE band extends from about 5700 to $\\rm 9100~\\AA$, with an estimated peak intensity near $\\rm 8300~\\AA$. Here, the ERE-to-scattered light band-integrated intensity ratio is about 0.7. At farther galactocentric distances, approaching the broad, bright HII region, the ERE band and peak intensity shift toward longer wavelengths, while the ERE band-integrated intensity diminishes. Finally, no ERE is detected within our spectral range and noise threshold at the inner edge of this HII region and at any other position at larger galactocentric distance. A secondary maximum of the ERE-to-scattered light band-integrated intensity ratio is found near the position of maximum opacity across the dust lane, associated with a secondary maximum of ongoing star formation activity. These variations reveal a complex dependence of the ERE on dust column density and interstellar radiation field. While the former quantity is derived from modeling, our data do not characterize the UV/optical flux shortward of $\\rm 5300~\\AA$. However, we are able to estimate the integrated (i.e. shortward of $\\rm 912~\\AA$) Lyman continuum photon rate and the characteristic effective temperatures of the OB-stars at its origin from the $\\rm H \\alpha$ ($\\rm \\lambda = 6563~\\AA$) line intensity and from the line intensity ratios $[NII](\\lambda~6583)/H \\alpha$ and $[SII](\\lambda \\lambda~6716+6731)/H \\alpha$. We interpret the ERE as originating from photoluminescence by nanometer-sized oxygen-passivated silicon particles, illuminated by UV/visible photons of the local radiation field. The phenomenology of the ERE in the Evil Eye is consistently interpreted through the recent model of the photophysics of the ERE carrier by Smith \\& Witt (2001), which attributes a key-role to the experimentally established recognition that photoionization quenches the luminescence of silicon nanoparticles. When examined within the context of ERE observations in the diffuse ISM of our Galaxy and in a variety of other dusty environments such as reflection nebulae and the Orion Nebula, we conclude that the ERE photon conversion efficiency in NGC4826 is as high as found elsewhere, but that the size of the actively luminescing nanoparticles in NGC4826 is about twice as large as those thought to exist in the diffuse ISM of our Galaxy." + }, + "0112/astro-ph0112297_arXiv.txt": { + "abstract": "We study the reionization of Helium {\\small II} by quasars using a numerical approach that combines 3D radiative transfer calculations with cosmological hydrodynamical simulations. Sources producing the ionizing radiation are selected according to an empirical quasar luminosity function and are assigned luminosities according to their intrinsic masses. The free parameters associated with this procedure are: (1) a universal source lifetime, (2) a minimum mass cutoff, (3) a minimum luminosity cutoff, (4) a solid angle specifying the extent to which radiation is beamed, and (5) a tail-end spectral index for the radiative energy distribution of the sources. We present models in which these parameters are varied and examine characteristics of the resultant reionization process that distinguish the various cases. In addition, we extract artificial spectra from the simulations and quantify statistical properties of the spectral features in each model. We find that the most important factor affecting the evolution of He {\\small II} reionization is the cumulative number of ionizing photons that are produced by the sources. Comparisons between He {\\small} II opacities measured observationally and those obtained by our analysis reveal that the available ranges in plausible values for the parameters provide enough leeway to provide a satisfactory match. However, one property common to all our calculations is that the epoch of Helium {\\small II} reionization must have occurred at a redshift between $3\\lesssim z \\lesssim 4$. If so, future observational programs will be able to directly trace the details of the ionization history of helium and probe the low density phase of the intergalactic medium during this phase of the evolution of the Universe. ", + "introduction": "Space-based ultraviolet telescopes have made it possible to observe the Ly$\\alpha$ transition of singly ionized helium (He {\\small II}) along lines of sight to high-redshift quasars. The discovery of the ``He {\\small II} Gunn-Peterson effect'' in Q0302-003 ($z=3.285$) by Jakobsen et al. (1994) marked the beginning of He {\\small II} Ly$\\alpha$ studies. He {\\small II} absorption in a second quasar, PKS 1935-692 ($z=3.18$), was identified by Tytler (1995; see also Jakobsen 1996), while Davidsen et al. (1996) measured the He {\\small II} opacity at a lower redshift of $z=2.72$ towards the quasar HS 1700+6416 using the Hopkins Ultraviolet Telescope (HUT). These papers established the presence of He {\\small II} absorption, albeit at low resolution, and indicated that the mean opacity increases with redshift. Later, Q0302-003 and HE 2347-4342 ($z=2.885$) were observed individually by Hogan, Anderson \\& Rugers (1997) and Reimers et al. (1997) using the Goddard High Resolution Spectrograph (GHRS). Subsequently, Heap et al. (2000) and Smette et al. (2000) reported new HST/STIS spectra for Q0302-003 and HE 2347-4342, respectively. These observations had sufficient resolution to begin to resolve some of the features in the He {\\small II} Ly$\\alpha$ forest and also allowed them to to cross-correlate the data with the absorbers in the gaps of the hydrogen (H {\\small I}) forest lines. In particular, the high quality of the STIS spectra reveals regions of high-opacity as well as ones of low opacity extending over several Mpc. Moreover, models of the spectra based on corresponding H {\\small I} Ly$\\alpha$ forests were presented to probe the hardness of the UV background, believed to emanate from the observed quasar population (e.g. Haardt \\& Madau 1996). Due to its relatively higher optical depth compared to H {\\small I}, He {\\small II} serves as a better probe of diffuse gas residing in voids between galaxies (e.g. Croft et al. 1997). Resolving He {\\small II} absorption features in quasar spectra can therefore reveal the presence of matter in large, low density regions, which, according to gravitational instability models of structure formation, harbor the bulk of the baryonic matter in the universe at high redshift (e.g. Dav\\'e et al. 2001; Croft et al. 2001). Interpreting observations of He {\\small II} absorption will require comparisons with detailed models derived from cosmological hydrodynamical simulations which incorporate radiative transfer effects responsible for the photoionization of helium. In particular, it is of interest to develop an understanding of how individual sources act collectively in the reionization process. In this paper we use the numerical approach described in Sokasian, Abel \\& Hernquist (2001, hereafter Paper I) to simulate the 3D reionization of He {\\small II} by quasars. In particular, we explore the parameter space associated with the characteristics of the sources and study how they influence global properties of the reionization process. Comparisons with observational results are also made possible by extracting synthetic spectra from the simulations. Here our aim is twofold: to develop an understanding of the sensitivity of the reionization process to source properties and to examine the predictions of the different models in light of recent observational results. ", + "conclusions": "The most important factor determining the reionization history of He {\\small II} is the cumulative number of ionizing photons that are produced by the sources. Given a specific form for the LF, ionizing emissivities are ultimately determined by intrinsic source characteristics such as their minimum luminosity and spectral energy distribution. In the analysis presented here, a numerical method designed to follow the inhomogeneous reionization of a cosmological volume by a set of point sources (quasars) was used to explore a small portion of the parameter space associated with these characteristics. Six models describing different sets of source properties were examined. Models 1-4 differ in source lifetimes and radiation beaming angles, but were all fixed to produce the roughly the same number of ionizing photons. More specifically, the tail-end spectral index of the SED in these models was set to $\\alpha_s=1.8$, as in the SED proposed by Madau, Haardt, \\& Rees (1999), while the minimum B-band luminosity at $z=0$ was set to $L_{min,o}=1.42 \\times 10^{9} \\ L_{B,\\odot}$ based upon an estimate from the results of Cheng et al. (1985). For comparison, models 5 and 6, were specifically designed to the produce substantially less ionizing emissivity while retaining plausible values for $\\alpha_s$ and $L_{min,o}$ (see Table 1). As expected, there is a sharp contrast between the results from the two groups. Specifically, the IGM in models 1-4 is ionized earlier and more thoroughly than in models 5 and 6, resulting in significantly lower mean opacities for the spectra (see Figure 3 and 11) at redshifts $z\\sim 3$. Based solely on the limited number of measured He {\\small II} opacities from observed quasar spectra, it appears that the ionizing emissivities in models 1-4 are too high, resulting in mean opacities lower than the observed values at all redshifts. If we accept our choice for $M_{min}$ as being reasonable, then the SED and/or $L_{min,o}$ need to be adjusted in order to better match the observations. In model 5, we picked a value of for $L_{min,o}$ that was a factor of 4 times higher than the value in models 1-4. Although this seems like a large difference from the original estimate of Cheng et al. (1985), a quick glance at Figure 1 of their paper reassures us that such a choice for the limiting luminosity (magnitude) for quasars is still plausible if we assume that the LF of Seyfert galaxies is similar to the LF of quasars. The tail-end spectral indices in models 5 and 6 were set to $\\alpha_s=2.0$ and $\\alpha_s=2.3$ respectively, which also represent plausible choices given the range of indices derived from observations of the EUV spectra of radio-quiet quasars at intermediate redshifts (Zheng et al. 1998). That model 5 provides the best match to observations of the mean optical depth bolsters the corresponding choices for $\\alpha_s$ and $L_{min,o}$. However, as model 6 makes clear, varying $\\alpha_s$ alone within a plausible range yields more than enough leeway to also match the observations while retaining the more widely quoted value for $L_{min,o}$ (see, e.g., Haardt \\& Madau 1996; Madau, Haardt, \\& Rees 1999; Bianchi et al. 2001). Furthermore, the results from model 3 suggest that the choice for source lifetimes may also affect the overall ionization state of the IGM, adding a further free parameter. However, uncertainties in the clumping, because of small scale power, and the gas temperatures do not allow us to clearly eliminate specific models using existing observational data. Therefore, we refrain from discussing best fits to the observations since we clearly have not conducted a comprehensive study of the associated parameter space. Rather, the main purpose of this paper is to provide a basic overview on how source characteristics influence global properties of the reionization process. Some specific insights from our analysis are as follows: 1. The systematic increase in the clumping factors associated with ionized regions towards low redshifts once most of the volume has already been ionized is consistent with the argument that in the post-reionization stage, the densest clouds and filaments become gradually ionized to higher degrees as the mean ionizing intensity increases. 2. The dispersion of the mean photoionization rates appears to be remarkably large and explicitly demonstrates the non-uniform nature of a radiation field dominated by local sources. Although dispersion levels do show signs of decreasing at the lowest redshifts when background intensities begin to exceed the intensities from local sources, the results demonstrate the importance of solving the radiation field locally during and before the epoch of reionization. The relatively large dispersion seen in models 2 and 3 can be attributed to the large number of sources which can lead to highly irregular radiation fields. 3. Measurements of the number of intersections as function of transmittance and the TPDF of Ly$\\alpha$ spectra offer insightful information regarding the ionization state of the IGM at specific times. In particular, both methods appear to be very sensitive to small differences in the underlying ionization fractions and as a result serve as useful tools to differentiate between spectra (models) that at first glance seem quite similar. 4. The lifetimes of the sources responsible for reionization may have an impact on the overall evolution of the process. Since sources tend to live in the densest regions of the IGM, the ionizing flux must first break through their dense surroundings in order to reach the IGM. Such an effect has been postulated for the spectrum of HS1700+6416 (Davidsen et al. 1996). One may therefore argue that the inability of model 3 to bring the IGM to high level of ionization relative to models 1, 2 and 4 might have to do with the fact that the relatively short lifetime requires a larger number of sources to be involved in reionization, effectively requiring more of the total ionizing flux to emanate from inside the densest structures. As a result, the IGM may be subject to less of the cumulative ionizing potential of the sources. One might wonder why model 2, which actually invokes even more sources than model 3 does not yield similar trends. In that case, the beaming of the radiation allows the sources to more easily \"break out\" of the dense absorbers which surround them thereby allowing more of the radiation to reach the diffuse IGM. The opposite case to model 3 occurs for model 4 where the longer lifetimes require fewer sources be invoked. In this case, by definition, there will be relatively fewer absorbers through which the sources have to shine their flux during the course of reionization. Consequently, the IGM in model 4 appears to be relatively more ionized than the other models. Our analysis reflects a first step to model radiative transfer effects and to develop new insights into open questions regarding the reionization of the universe. In particular, the method used here can be modified to study the reionization of H {\\small I}, offering further opportunities for comparisons with observations, especially in light of recent evidence for the reionization of H {\\small I} at $z\\sim6$ (Becker et al. 2001). Furthermore, analysis of both H {\\small I} and He {\\small II} Ly$\\alpha$ opacities from simulated spectra and observational results can constrain the contribution of starburst galaxies to the UV ionizing background (see, e.g., Bianchi et al. 2001; Kriss et al. 2001). We also envision supplementing our method to include the effect of heating from photoionization to explore the possibility of searching for fluctuations in the IGM temperature using the Ly$\\alpha$ forest (see, e.g, Zaldarriaga 2001)." + }, + "0112/astro-ph0112542_arXiv.txt": { + "abstract": "The detection of two similar periodicities (3001 and 3028~s) in the light curve of V1405~Aql, a low mass X-ray Binary (LMXRB), has attracted the attention of many observers. Two basic competing models have been offered for this system. According to the first, V1405~Aql is a triple system. The second model invokes the presence of an accretion disc that precesses in the apsidal plane, suggesting that the shorter period is the orbital period while the longer is a positive superhump. The debate on the nature of V1405~Aql has been continued until very recently. Re-examination of previously published X-ray data reveals an additional periodicity of 2979~s, which is naturally interpreted as a negative superhump. The recently found 4.8-d period is consequently understood as the nodal precession of the disc. This is the first firm detection of negative superhumps and nodal precession in a LMXRB. Our results thus confirm the classification of V1405~Aql as a permanent superhump system. The 14-year argument on the nature of this intriguing object has thus finally come to an end. We find that the ratio between the negative superhump deficit (over the orbital period) and the positive superhump excess is a function of orbital period in systems that show both types of superhumps. This relation presents some challenge to theory as it fits binaries with different components. We propose that a thickening in the disc rim, which causes increased occultation of the X-ray source, is the mechanism responsible for both types of superhumps in LMXRBs. However, the positive signal is related only to the pronounced dips in the light curve, where the point-like central source is covered up, whereas the morphology of the negative superhump signal appears quite smooth, implying obscuration of a larger X-ray emitting region, possibly the inner accretion disc or a corona. According to our model superhumps (both in the X-ray and optical regimes) are permitted in high inclination LMXRBs contrary to Haswell et al. (2001) prediction. ", + "introduction": "\\subsection{Permanent superhumps} Permanent superhumps have been observed so far in about 20 cataclysmic variables (CVs) (Patterson 1999). These show superhumps (quasi-periodicities shifted by a few percent from their orbital periods) in their optical light curves during normal brightness state. In contrast, SU~UMa systems (see Warner 1995 for a review of SU~UMa systems and CVs in general) have superhumps only during their bright dwarf nova outbursts (superoutbursts). Permanent superhumps can either be a few percent longer than the orbital periods and they are called `$\\bf positive$ $\\bf superhumps$', or shorter -- `$\\bf negative$ $\\bf superhumps$'. The positive superhump is explained as the beat between the binary motion and the precession of an accretion disc in the apsidal plane. Similarly, the negative superhump is understood as the beat between the orbital period and the nodal precession of the disc (Patterson 1999). According to theory, superhumps are developed when the accretion disc extends beyond the 3:1 resonance radius and becomes elliptical (Whitehurst \\& King 1991). Osaki (1996) suggested that permanent superhump systems differ from other subclasses of non-magnetic CVs in their relatively short orbital periods and high mass-transfer rates, resulting in accretion discs that are thermally stable but tidally unstable. Retter \\& Naylor (2000) gave observational support to this idea. Simulations suggest that superhumps can only occur in binary systems with small mass ratios -- $q=M_{donor}/M_{compact}\\la0.33$ (Whitehurst 1988; Whitehurst \\& King 1991; Murray 2000). Although this condition is easily met in short orbital period LMXRBs, whose primaries are compact massive objects, superhumps have been only seen in a few LMXRBs in outburst (e.g. O'Donoghue \\& Charles 1996). Here we report the detection of permanent superhumps in a persistent LMXRB. \\subsection{V1405~Aql (X1916-053 / 4U 1915-05)} V1405~Aql was first discovered as an X-ray source showing type-I bursts (Becker et al. 1977; Doxsey et al. 1977) suggesting that its primary star is a neutron star. White \\& Swank (1982) and Walter et al. (1982) independently found $\\sim$3000-s periodic dips in its X-ray light curve. The dips are believed to occur due to obscuration of the disc edge (probably where the accretion stream impacts the disc). Such a short orbital period implies a low mass hydrogen-deficient secondary star (Nelson, Rappaport \\& Joss 1986). The 21-mag optical counterpart of V1405~Aql was identified by Schmidtke (1988) and Grindlay et al. (1988). They also reported a detection of an optical periodicity, about one percent longer than the X-ray period. The difference between the X-ray and optical periods was confirmed by further extensive observations of V1405~Aql (Smale et al. 1988; 1989; 1992; Callanan 1993; Callanan, Grindlay \\& Cool 1995; Yoshida et al. 1995; Church et al. 1997; 1998; Ko et al. 1999; Morley et al. 1999; Chou, Grindlay \\& Bloser 2001; Homer et al. 2001). The various models offered so far for the periodicities found in V1405~Aql can be grouped into two basic models -- a triple system (Grindlay 1986; Grindlay et al. 1988) and superhumps (Schmidtke 1988; White 1989). According to the first model, which was motivated by the detection of a possible 199-d periodicity in the X-ray light curve of V1405~Aql (Priedhorsky \\& Terrell 1984), the longer 3028-s period is the binary inner orbital period, while the shorter 3001-s period is the beat between the binary period and the $\\sim$4-d orbital period of a third companion. The 199-d period is explained by the eccentricity of the inner binary orbit (Chou et al. 2001). The second model suggests that the 3001-s period is the binary period and that the 3028-s period is a positive superhump. We note that both models predict the presence of a beat periodicity of $\\sim$4 d. Indeed evidence for a period of $\\sim$3.9~d has been found in the X-ray light curve of V1405~Aql (Chou et al. 2001; see also Homer et al. 2001). The debate on the nature on V1405~Aql has continued (Chou et al. 2001; Homer et al. 2001; Haswell et al. 2001), although some preference was given to the superhump scenario. One of the arguments against the triple system model is that the 199-d period was not confirmed by further observations, although a second possible evidence for this periodicity was found by Smale \\& Lochner (1992). On the other hand, an argument against the superhump model is that the 3028-s period is quite stable while superhump periods are unstable (Callanan et al. 1993; 1995; Chou et al. 2001; Homer et al. 2001). There is a strong observational link between positive and negative superhumps in CVs as light curves of several systems show both (Patterson 1999; Arenas et al. 2000; Retter et al. 2002b). In addition, Patterson (1999) found that period deficits in negative superhumps are about half period excesses in positive superhumps: $\\epsilon_{-}\\approx-0.5\\epsilon_{+}$, where $\\epsilon=(P_{superhump}-P_{orbital})/P_{orbital}$. Thus, we decided to look in available photometric data on V1405~Aql for negative superhumps, which would be predicted to have a period near 2986~s. Indeed we have found a new periodicity in the X-ray light curve of V1405~Aql. These findings were reported by Retter, Chou \\& Bedding (2002a), and are presented below. We also present a new relation for systems that have both types of superhumps, and discuss the mechanisms for superhumps in the X-ray and optical regimes. ", + "conclusions": "\\subsection{The new period} A period of 2979~s has been found in the X-ray light curve of V1405~Aql in addition to the other known periods. This feature was actually noticed by Chou et al. (2001) who mistakenly identified it with a 3.9-d sideband of the 3001-s period. It is, however, $\\sim5\\sigma$ away from the expected sideband at 2974.2 s and we have shown above that it is not an alias. We also note that a period of 2984.6$\\pm$6.8 s was detected in the Ginga data of this object (Smale et al. 1989). As it is inconsistent with either the 3001 and 3028-s periods, it probably represents the new periodicity as well. The 2979-s period is shorter than the 3001-s period by about 0.7 percent. Assuming that the 3001-s period is the orbital period and that the 3028-s period is a positive superhump, the new period is naturally explained as a negative superhump. Our suggestion implies a nodal precession of $\\sim$4.8 d. Indeed Chou et al. (2001) found that the phase jitter of the X-ray dips in the 1996 May data is modulated with a period of 4.86 d (as well as with the $\\sim$3.9-d period). Homer et al. (2001) reached a similar conclusion from a different dataset and found a period of 4.74$\\pm$0.05 d. Thus, the superhump model can explain this periodicity as well. The 2979-s period in the X-ray light curve of V1405~Aql and the beat period at $\\sim$4.8 d cannot be explained by the triple system model. This scenario was motivated by the possible detection of the 199-d period that has not been confirmed by further observations (Section 1.2). Homer et al. (2001) suggested a few other possible periodicities ($\\sim$8.6, $\\sim$12, $\\sim$23, $\\sim$83 d) in the X-ray light curve of V1405~Aql. It is very likely that V1405~Aql has random variations of the order of a few tens of day, but not a coherent long-term periodicity. We thus believe that the triple system model for V1405~Aql is finally ruled out. Table~1 lists the periods observed in V1405~Aql and their interpretations according to the superhump model. The presence of an extensive number of periodicities can explain many peculiarities in the light curve of the system that have been previously discussed by many authors. \\begin{table} \\caption{The periods in V1405~Aql} \\begin{tabular}{@{}clll@{}} Number & Period & Nature & Reference \\\\ 1 & 2.8 ms & spin & Boirin et al. (2000) \\\\ & 3.7 ms & & Galloway et al. (2001) \\\\ & & & \\\\ 2 & 2979.3(1.1) s & negative & this work \\\\ & 2984.6(6.8) s & superhump & Smale et al. (1989) \\\\ & & & \\\\ 3 & 3000.6508(9) s & orbital & Chou et al. (2001) \\\\ & 3000.6(2) s & & Homer et al. (2001) \\\\ & & & \\\\ 4 & 3027.5510(52) s & positive & Chou et al. (2001) \\\\ & 3027.555(2) s & superhump & Homer et al. (2001) \\\\ & & & \\\\ 5 & 3.9087(8) d & apsidal & Chou et al. (2001) \\\\ & & precession & \\\\ & & & \\\\ 6 & 4.86 d & nodal & Chou et al. (2001) \\\\ & 4.74(5) d & precession & Homer et al. (2001) \\\\ \\end{tabular} \\end{table} \\subsection{A new relation for superhumps} Patterson (1999) proposed that the negative superhump deficit is about half the positive superhump excess (Section~1.2). As the 3028-s period is about 0.9 percent longer than the 3001-s period, the corresponding ratio in V1405~Aql (0.79) is somewhat larger than this. We have checked the connection between the deficit and excess of negative and positive superhumps in CVs and found a new relation. Table~2 presents the data on the periods in systems showing both types of superhumps. In Fig.~5 we show this ratio for these systems, and we see a clear trend as a function of orbital period. Our result for V1405~Aql fits this trend very well. \\begin{table*} \\begin{minipage}{300 mm} \\caption{Properties of systems that have the two kinds of superhumps} \\begin{tabular}{@{}lccccccc@{}} Object & Orbital & Positive & Positive Superhump & Negative & Negative Superhump & $\\phi=\\epsilon_{-}/\\epsilon_{+}$ & Ref. \\\\ & Period & Superhump & Excess ($\\epsilon_{+}$) & Superhump & Deficit ($\\epsilon_{-}$) & & \\\\ & [d] & [d] & & [d] & & & \\\\ AM~CVn&0.011906623(3)&0.0121666(13)& 0.02183(12) &0.01170613(35)& -0.016839(30) & -0.7714(57) & 1,2,3 \\\\ V1405~Aql&0.034729754(11)&0.035041099(60)& 0.0089648(20) &0.034483(13)& -0.0071(7) & -0.79(8) & 4,5 \\\\ V503~Cyg &0.0777(2)& 0.08104(7)& 0.0430(27) &0.07597(18)& -0.022(5) & -0.51(16) & 2,3,6 \\\\ V1974~Cyg&0.0812585(5)&0.08506(11)& 0.0468(14) & 0.07911(5)& -0.02644(62) & -0.565(31) & 2,3,7,8 \\\\ TT~Ari &0.1375511(2)& 0.1492(1)& 0.0847(7) & 0.1329(3) & -0.034(3) & -0.401(40) & 2,3,9 \\\\ V603~Aql& 0.1381(1) & 0.1460(7)& 0.0572(51) & 0.1343(3) & -0.028(2) & -0.490(86) & 2,3,10 \\\\ TV~Col & 0.22860(1) &0.2639(35)& 0.154(15) & 0.2160(5) & -0.0551(23) & -0.357(55) & 3,11 \\\\ \\end{tabular} \\end{minipage} $^1$Skillman et al. (1999); $^2$Patterson (1998); $^3$Patterson (1999); $^4$Chou et al. (2001); $^5$This work; $^6$Harvey et al. (1995); $^7$Retter, Leibowitz \\& Ofek (1997); $^8$Skillman et al. (1997); $^9$Skillman et al. (1998); $^{10}$Patterson et al. (1997); $^{11}$Retter et al. 2002b; \\vspace{0.25cm} Notes to Table~2: 1. V1159~Ori (Patterson et al. 1995), V592~Cas (Taylor et al. 1998), PX~And and BH~Lyn (Patterson 1999) were rejected from this sample as their superhump periods are still uncertain (Patterson, personal communication). 2. ER~UMa (Gao et al. 1999) was not considered as the detection of negative superhumps does not seem secure. \\end{table*} \\begin{figure} \\centerline{\\epsfxsize=3.0in\\epsfbox{fig5.eps}} \\caption{The relation between the orbital period and the ratio between the negative superhump deficit and the positive superhump excess in systems that have both types of superhumps.} \\end{figure} \\subsection{Implications} Our results strengthen the observational link between positive and negative superhumps, and supports the idea that the two types of superhumps have a similar physical origin, namely a precessing accretion disc. Retter et al. (2002b) speculated that every permanent superhump system may have both kinds of superhumps. The data on V1405~Aql support this idea. Wood, Montgomery \\& Simpson (2000) showed that a tilted accretion disc can explain the presence of negative superhumps, however, it is still unclear what physical force drives the disc to precess in the nodal plane (Murray \\& Armitage 1998). Wynn (personal communication) suggested that a strong magnetic field of the primary white dwarf can cause the disc to precess in the nodal plane. If this idea is confirmed, negative superhumps are expected in every permanent superhump LMXRB, as their primaries have magnetic fields typically much stronger than in CVs. Montgomery (2001) and Montgomery et al. (in preparation) developed analytic expressions for the two types of superhumps. The theoretical values for the ratio $\\phi$=$\\epsilon_{-}/\\epsilon_{+}$ mentioned above, were found to be consistent with the observations of CVs. Montgomery et al.'s results can also explain the relation shown in Fig.~5 for CVs (see their Fig.~6). Using their equation 9 \\& Fig.~4 (but baring in mind the difference in primary mass) and the observed $\\phi$ parameter in V1405~Aql (Table~2), we can derive a weak upper limit on its mass ratio of $q\\la0.08$. For a primary neutron star at the Chandrasekhar limit, this implies an upper limit on the mass of the secondary -- $M_{2}\\la0.12M_{\\odot}$, which is presumably a helium white dwarf. The fact that binary systems with very different configurations (CVs -- a hydrogen-rich white dwarf and a red companion; AM~CVn -- two degenerate helium white dwarfs and V1405~Aql -- a neutron star and a helium white dwarf) obey the relation presented in Fig.~5 is quite surprising. It might be a coincidence, or it may suggest that $\\phi$ indeed depends on the orbital period and not on the mass ratio. Further determinations of this parameter in more permanent superhump systems are required to study this issue. \\subsection{The mechanism for negative superhumps} The mechanism which places the negative superhump signal into the X-ray light curve is unclear, but changing vertical disc structure seems a promising candidate. It is clear that positive superhumps coincide with increased disc thickness (Billington et al. 1996), the evidence being consistent with the idea that the area of the disc involved in producing the superhump light thickens at the time of the increased dissipation. In the nodal precession model, Wood et al. (2000) showed how the passage of the secondary star past each of the two halves of the disc out of the midplane increases the dissipation in that half. They also required that the side of the disc closest to the orbital plane is the most disrupted. If, by analogy with positive superhumps, this disruption increases the vertical extent of the disc, then it will increase the obscuration of the X-ray source. As the disc precesses, the timing of the obscuration will change with phase, producing the required modulation at the negative superhump period. The lack of structure and large range in phase of the negative superhump modulation implies that both the disc structure and occulted region are large, the latter presumably being the inner accretion disc or a corona. This scenario can also explain why negative superhump signals have not been observed so far in the X-ray light curves of low inclination systems. \\subsection{The mechanism for positive superhumps} Haswell et al. (2001) pointed out that the dissipation of energy in the disc, which is believed to be responsible for the positive superhump mechanism in CVs in the optical, cannot be applied to LMXRBs. They suggested that instead the disc area is changing with the superhump period, and thus predicted that superhumps would appear mainly in low inclination systems. Figs.~1 \\& 2 show that the X-ray data is modulated with the positive superhump (f$_{2}$). However, when the dips are rejected from the light curve, this peak disappears. Thus in V1405 Aql only the dips show the positive superhump, presumably since their amplitude and / or phase vary with the apsidal precession period. Note that the phase jitter of the dips is modulated with the 3.9-d period (Chou et al. 2001). Therefore, a simple explanation of this behaviour is that thickening of the disc rim, which causes an increase obscuration of the X-ray source, forms the positive superhump. This is consistent with our suggestion for the mechanism for the negative superhump in the X-ray (previous section). The disc thickening would allow the disc to intercept more of the irradiating flux, thus producing optical superhump light seen at all inclinations in a similar way to the Haswell et al. area effect. In addition, though, it would introduce dips into the X-ray and perhaps optical orbital light curves, allowing the superhump to be visible in high inclination systems, in contrast to the prediction of Haswell et al." + }, + "0112/astro-ph0112068_arXiv.txt": { + "abstract": "{From the South Pole, microthermal turbulence within a narrow surface boundary layer some 200\\,m thick provides the dominant contribution to the astronomical seeing. We present results for the seeing at a wavelength of 2.4$\\mu$m. The narrow turbulence layer above the site, confined close to the surface, provides greatly superior conditions for adaptive optics correction than do temperate latitude sites. An analysis of the available meteorological data for the Antarctic plateau suggests that sites on its summit, such as Domes A and C, probably experience significantly better boundary layer seeing than does the South Pole. In addition, the inversion layers may be significantly narrower, lending the sites even further to adaptive optics correction than does the Pole. ", + "introduction": "The effect of atmospheric turbulence on astronomical image quality, or ``seeing'', has been studied at the South Pole through measurement of the microthermal fluctuations associated with the turbulence (Marks et al.\\ 1996---Paper I, Marks et al.\\ 1999---Paper II). The seeing was found to be dominated by the contributions from a narrow but turbulent boundary layer, with a minimal contribution from the free atmosphere above it. This suggests that, if the effects of the surface boundary layer over the Antarctic plateau can be mitigated, superb seeing conditions might be obtained. Since the boundary layer seeing is strongly influenced by the inversion wind, it is possible that exceptionally good seeing may occur at the surface from other locations on the plateau away from the South Pole, where this wind is reduced. Such sites may be the high points of the Antarctic plateau---Domes A and C---where the inversion wind is almost non-existent. Little direct evidence is available on seeing conditions for these sites. Some clues, however, are available from meteorological records. We examine these in this paper, and discuss the implications for future astronomical observatories in Antarctica. While the South Pole is convenient location, in the sense that it is populated all year-round, it is generally agreed that there are other sites where the observing conditions are potentially much better, in terms of atmospheric transmission, humidity and weather, in addition to the seeing. The South Pole lies some 1,000\\,km away from Dome~A, the highest point on the plateau at about 4,200\\,m. Much of the boundary layer turbulence at the Pole is associated with the inversion winds which consist of cold air rolling gradually down the slope from the highest regions of the plateau, picking up speed as they go, before finally turning into the violent katabatic winds that are such a well-known feature of the weather along the Antarctic coastline. It is quite possible that the boundary layer seeing at Dome A is much lower than at the South Pole since, although the temperature inversion is still present, the calm winds close to the surface mean that the mechanical mixing of the different temperature layers is minimised. Since Dome A still remains almost totally inaccessible it is fortunate that there are other sites for which similar comments apply. Potential candidates include Vostok and Dome~C\\@. Dome~C has the further advantage that the infrastructure will soon be in place at this site to support a large observatory, with construction of the year-round Franco-Italian {\\sl Concordia Station}, expected to be completed by 2004. Obviously, direct measurements need to be made at these sites; however the sparse information available at this stage does allow some general comments to be made about the likely seeing conditions at the higher plateau sites, given the results from the South Pole. ", + "conclusions": "To summarise our results, the microthermal turbulence at the South Pole is concentrated much closer to the surface than at the best mid-latitude sites. So, while the integrated seeing from surface level is poor, the free atmosphere is quieter than at these other sites. This has the result that the seeing contribution above about 200~m (the depth of the South Pole boundary layer) is significantly less than at any temperate latitude site. The available evidence also suggests that the high-altitude sites such as Domes~A and C probably experience better boundary layer seeing than the South Pole. Given that this is by far the dominant source of image degradation at the Pole, this is a very important result. The conclusion is based on the absence of the inversion winds that generate much of the wind shear responsible for the mechanical turbulence. In addition, the inversion layers themselves may be significantly narrower at the higher sites. The narrow turbulent region at the South Pole, relative to mid-latitude sites, greatly increases the characteristic angles over which image correction techniques can be applied. The higher sites may also, therefore, be significantly better in terms of their potential for adaptive optics." + }, + "0112/astro-ph0112224_arXiv.txt": { + "abstract": "{ Spectroscopic observations covering the wavelength range 3600--4600\\AA\\ are presented for a large flare on the late type M dwarf AT Mic (dM4.5e). A procedure to estimate the physical parameters of the flaring plasma has been used which assumes a simplified slab model of the flare based on a comparison of observed and computed Balmer decrements. With this procedure we have determined the electron density, electron temperature, optical thickness and temperature of the underlying source for the impulsive and gradual phases of the flare. The magnitude and duration of the flare allows us to trace the physical parameters of the response of the lower atmosphere. In order to check our derived values we have compared them with other methods. In addition, we have also applied our procedure to a stellar and a solar flare for which parameters have been obtained using other techniques. ", + "introduction": "Stellar flares are events where a large amount of energy is released in a short interval of time, radiating at almost all frequencies in the electromagnetic spectrum. Flares are believed to result from the release of magnetic energy stored in the corona through reconnection \\citep[see reviews by][]{Mirzoyan84,Butler91,Haisch91,Garcia-Alvarez00}. Many types of cool stars produce flares \\citep{Pettersen89}, sometimes at levels several orders of magnitude more energetic than their solar counterparts. The exact mechanism(s) leading to the energy release and subsequent excitation of various emission features remains poorly understood. In the dMe stars (or UV Cet type stars) optical flares are a common phenomenon. In more luminous stars, flares are usually only detected through UV or X-ray observations \\citep{Doyle89b}, although optical flares have been detected in young early K dwarfs like LQ Hya \\citep{Montes99}. One would like to trace all the energetic processes in a flare to some common origin likely to be localized at some unresolvably fine scale in the magnetic field. {The largest solar flares observed involve energies of $\\sim10^{32}\\ \\rm{erg}$ \\citep{Gershberg89}, while large flares on dMe stars can be one order of magnitude larger \\citep{Doyle90,Byrne90}}. Even more energetic flares occur on the RS CVn binary systems, the total flare energy in the largest of this type of systems may exceed E$\\sim 10^{38}\\ \\rm{erg}$ \\citep{Doyle92,Foing94}. Such a change in the star's radiation field modifies drastically the atmospheric properties over large areas, from photospheric to coronal layers. {Models by \\citet{Houdebine92} indicate that heating may propagate down to low photospheric levels, with densities higher than $10^{16}\\ \\rm{cm}^{-3}$, although electrons with energies in the $\\rm{MeV}$ range would be required to attain such depth.} As regards the derivation of physical parameters such as electron density and temperature, various methods have {been} used. For example, \\citet{Katsova90} published an analysis of flare Balmer decrements. The broadening and merging of higher Balmer lines, dominated by the Stark effect, have also {been} used to estimate electron densities in the chromosphere \\citep{Donati-Falchi85}. Measurements of the broadening of the lower Balmer lines, which are less affected by the Stark effect, together with line shifts, provide information on the large scale motions during flares. {Steep decrements are evidence for electron densities between $10^{8}$ and $10^{12}\\ \\rm{cm}^{-3}$, while a shallow Balmer decrement indicates densities larger than $10^{13}\\ \\rm{cm}^{-3}$ in the flaring M dwarf atmosphere.} \\citet{Jevremovic98} and \\citet{Jevremovic99} developed a procedure to fit the Balmer decrement based on the solution of the radiative transfer equation using the escape probability computing technique of \\citet{Drake80} and \\citet{DrakeUlrich80} and the direct search method of \\citet{Torczon91,Torczon92}. In this paper we present the results of a large flare observed on AT Mic (Gliese 799B, V=10.25, B=11.83 $ 20^h41^m51^s, -32^o26'07'' $, Equinox 2000). It is a visual binary star with a separation of 4.0 arcsec \\citet{Wilson78}, approximately 10.2 pc away, with both components of spectral type dM4.5e and subject to flaring \\citep{Joy49}. \\citet{Kunkel70} determined the flare incidence at 2.8 flares per hour. From its position in the (Mv, R-I) diagram, \\citet{Kunkel73} deduced that AT Mic lies slightly above the main sequence, pointing to its probable membership of the young disc population. \\citet{Nelson86} reported 29 U-band/B-band flares in 39.3 hours, only one of which showed microwave flaring at 5 \\rm{GHz}. They also reported two microwave bursts without any optical counterpart. \\citet{Kundu87} found that both components of the AT Mic system were active and variable at 6 and 20 $\\rm{cm}$ wavelengths. The peak emission levels were detected in the southern component (Gliese 799B). X-ray and ultraviolet emission from this system has been observed several times over the last ten years or so \\citep{Linsky82,Pallavicini90}. In $\\S$2 we present the {current} dataset plus a discussion of the data analysis while $\\S$3 gives the results. In $\\S$4, we analyse the evolution of the main physical parameters during the flare using the code developed by \\citet{Jevremovic98} and apply the procedure to a stellar and a solar flare for which parameters have been obtained using other techniques. The conclusions are given in $\\S$5. ", + "conclusions": "We observed a major flare on AT Mic in the 3600--4500\\AA\\ wavelength region which allowed us to compute a detailed flare energy budget in the optical spectral region. The physical parameters obtained are consistent with previously derived values for stellar flares. We have compared the results from our method for two flares previously analysed, namely, one on YZ CMi and one on the Sun, obtaining reasonable agreement. We have obtained the first detailed trace of physical parameters during a stellar flare. We have also shown that our method could provide a suitable model for the study of flares on stars of different spectral type (G-M). The total radiated energy released in the region 3600-4500\\AA\\ during the flare was $\\sim4\\,10^{33}\\ \\rm{erg}$ and $\\Delta U\\sim\\,4\\, \\rm{mag}$. The energy radiated by this flare is comparable with the strongest solar ones but is of medium strength compared to flares on other dMe stars. Further stellar observations with high time and high spectral resolution should be obtained. Also, further solar observations should be obtained in order to check the method with other examples, particularly, the flare area and thickness." + }, + "0112/astro-ph0112538_arXiv.txt": { + "abstract": " ", + "introduction": "In the past years and decades, several models of nova shells have been presented in the literature. Often they were adapted to describe the state and evolution of specific objects, and often remarkable agreement between model and observation was achieved. Nevertheless it should be kept in mind that a nova shell is a rapidly evolving object, and its properties change significantly with time. Furthermore, a plethora of different types of novae are observed, which is accompanied by an amazing variety of nova shells of various morphologies and physical properties in different stages of temporal development. Although studies of nova shells have been carried out since the first bright nova of the 20th century, GK Persei in 1901, most of these studies were carried out in a qualitative way. This approach permitted the calculation of nebular expansion parallaxes and the morphological study of shells. Since the shells were usually faint, and the observations were carried out with photographic plates, hardly any quantitative results are available. Only in the first phases of the outburst, when the shells document themselves in the form of emission lines, were the line fluxes estimated and derived for a few cases, notably by Payne-Gaposchkin and collaborators. Replacement of the photographic plate by digital receivers has facilitated the task of studying the evolution of nova remnants, both spectroscopically and by means of direct imaging through narrow-band filters. In fact, quite a number of studies have even been carried out for extragalactic novae, where \\Hza{}-images can more easily detect the objects above the stellar background (see, e.g. Ciardullo et al. 1987). In this paper, we report on the results of a recent imaging survey of nova remnants, carried out at the Kitt Peak and ESO La Silla observatories. We also use a hitherto unpublished survey of nova shells carried out in 1984 at Calar Alto, and the images from the \\HST archive. Furthermore, we have collected and homogenized the existing quantitative record of nova shell observations. Because the survey attempted to cover as many objects in as many evolutionary stages as possible, hardly any detailed information on a given object, or any detailed modelling of shells will be given (i.e. the distribution of line flux between various specific parts of a nova shell). We rather attempt to describe the ``average'' or global evolutionary track of a nova shell, in order to derive expected values for faint shells of ancient novae. A theoretical interpretation of the observed behavior will be the subject of a forthcoming paper (Duerbeck \\& Downes 2002). Section 2 describes our observations and reductions. Section 3 briefly describes the classification of novae according to speed class, which is the base for merging our shell luminosity data into groups. Section 4 gives the derivation of global trends in luminosity evolution for the lines \\Hza{}, \\Hzb{} and \\OIII{} in novae of different speed classes (including, besides classical novae, recurrent ones). Section 5 summarizes our results. ", + "conclusions": "We have collected about 1200 available line fluxes of 96 classical and recurrent novae of various speed classes and have studied the evolution of the luminosity in the \\Hza, \\Hzb, and \\OIII line as a function of time after outburst. In general, novae of a given speed class follow similar patterns, so that functional relations for the average evolution of line luminosity could be derived. General trends for novae of various speed classes are shown in Figs. 15, 16 and 17. A few novae turn out to be unusual: GK Per and T Pyx, which interact with circumstellar material, V838 Her and V4160 Sgr, which have unusually small mass ejection, and GQ Mus, an X-ray emitting classical nova. \\begin{figure} \\centerline{\\psfig{figure=finchaall.eps,width=9cm,angle=270}} \\caption{Averaged time dependence of \\Ha emission in novae of various speed classes: very fast (blue), fast (green), moderately fast (yellow), slow and very slow (red), and recurrent novae with giant companions (black). At late times, decay occurs fastest in very fast novae, somewhat more slowly in slow novae, and slowest in fast and moderately fast novae. Contrary to the behaviour in classical novae, the decay of \\Ha emission in recurrent novae sets in immediately after maximum, and continues without any break in slope.} \\end{figure} \\begin{figure} \\centerline{\\psfig{figure=finchball.eps,width=9cm,angle=270}} \\caption{Averaged time dependence of \\Hb emission in novae of various speed classes: very fast (blue), fast (green), moderately fast (yellow), slow and very slow (red), and recurrent novae with giant companions (black). Note that the behaviour of \\Hb in novae of different speed classes is quite similar to that shown in \\Hza. Contrary to the behaviour in classical novae, the decay of \\Hb emission in recurrent novae sets in immediately after maximum, and continues without any break in slope.} \\end{figure} \\begin{figure} \\centerline{\\psfig{figure=finco3all.eps,width=9cm,angle=270}} \\caption{Averaged time dependence of \\OIII emission in novae of various speed classes: very fast (blue), fast (green), moderately fast (yellow), slow and very slow (red), and recurrent novae with giant companions (black). At late times, decay occurs fastest in very fast novae, almost as fast in moderately fast and fast novae, and slowest in slow novae, i.e. shells of slow novae show noticeable \\OIII emission also at late times. \\OIII emission in recurrent novae appears to be ``on'' for some time and afterwards disappears quickly. } \\end{figure} A general discussion of the material presented here in the framework of nova properties and shell evolution will be presented elsewhere (Duerbeck \\& Downes 2002). The data of Tables 1 -- 4 are also included as ascii files in {\\tt table1.dat, table2.dat, table3.dat} and {\\tt table4.dat}." + }, + "0112/astro-ph0112362_arXiv.txt": { + "abstract": "We have undertaken a program of high-resolution imaging of high-redshift radio galaxies (HzRGs) using adaptive optics on the Canada-France-Hawaii Telescope. We report on deep imaging in $J$, $H$, and $K$ bands of 6 HzRGs in the redshift range 1.1 to 3.8. At these redshifts, near-infrared bandpasses sample the rest-frame visible galaxian light. The radio galaxy is resolved in all the fields and is generally elongated along the axis of the radio lobes. These images are compared to archival Hubble Space Telescope Wide-Field Planetary Camera 2 optical observations of the same fields and show the HzRG morphology in rest-frame ultraviolet and visible light is generally very similar: a string of bright compact knots. Furthermore, this sample - although very small - suggests the colors of the knots are consistent with light from young stellar populations. If true, a plausible explanation is that these objects are being assembled by mergers at high redshift. ", + "introduction": "The range of redshifts of the known radio galaxies (RGs) spans from $z=0.06$ for \\objectname{Cygnus A} to well above $z=3$. For example, \\objectname{4C +41.17} at $z=3.79$, \\objectname{6C 0140+326} at $z=4.41$, and \\objectname{TN J0924-2201} at $z=5.19$. It is well established that the most powerful low redshift radio sources are associated with giant elliptical galaxies, often the brightest cluster galaxies. The optical morphology of powerful RGs with redshifts greater than $\\sim 0.6$ is dramatically different from that of those of lower redshift. These high-redshift radio galaxies (HzRGs) are clumpy and very irregular and tend to be elongated along the axis of their radio lobes, a characteristic known as the `alignment effect' \\citep{Chambers1987, McCarthy1987}. The best optical imaging of HzRGs has been obtained with the Hubble Space Telescope (HST) Wide-Field Planetary Camera 2 (WFPC2). These data include $\\sim 100$ objects mostly from the 3C and 4C catalogs and have revealed very complex rest-frame UV morphologies down to the resolution limit of $\\approx 0.1${\\arcsec} \\citep{Longair1995, Best1996, Best1997, Chambers1996, McCarthy1997, Pentericci1999}. These images reveal the immediate neighbourhood of the AGN host to be strings of knots with separations on the order of a few kpc aligned along the radio axis. These are embedded in a diffuse emission region with projected scales of $\\sim 50$ kpc. Typically, several faint field objects within $\\sim 100$ kpc projected radius from the host are also found. A still unanswered question about the stellar population of HzRGs is whether it is young or old. Since the highest redshifts for RGs are in the 3 - 5 range, one might assume that they are very young systems since, for a universe with $q_0=0.5$, the look-back time is about 90\\% of the age of the universe. A popular theory presented by \\citet{Lilly1984} \\citep{Lilly1988, Lilly1989} is that most of the star formation in these systems took place in an initial burst at $z_{\\rm formation}\\sim 5-10$ and that minor star-formation episodes afterwords produce the dramatic morphologies in rest-frame UV light. This `old galaxy $+$ burst' model is reasonably successful in matching the $K-z$ relation for RGs but it is not the only way to account for the rest-frame UV light. The high incidence of double and multiple component galaxies with separations of a few kpc suggests that these may be in the process of merging in the direction defined by the radio axis \\citep{Djorgovski1987, West1991, West1994}. Star formation could be induced by the shocking of the intergalactic medium by jets along the radio axis \\citep{Chambers1987, McCarthy1987, DeYoung1989, Rees1989, Begelman1989, Chambers1990, Daly1990}. The high linear polarization of some galaxies suggests that scattering of light from the active galactic nucleus (AGN) by electrons or dust may be the dominant source of emission \\citep{diSeregoAlighieri1989, Tadhunter1992, diSeregoAlighieri1993, Januzzi1995, Cimatti1996, Cimatti1997}. Also, nebular thermal continuum associated with ionized gas is an important contributor for some objects \\citep{Dickson1995, Stockton1996}. The rest-frame UV morphologies of HzRGs may be dramatic but in order to study galaxy evolution in HzRGs one should follow the mature stellar populations. If these populations are mapping out the structure of the `true' galaxy one might well ask how this structure changes in time. Do these galaxies appear as elliptical galaxies from early epochs to the present or, if not, how do they evolve? Deep, high-resolution NIR imaging allows the study of HzRGs to be extended to the most distant ones known. It permits the investigation of the rest-frame-visible properties of RGs over a large range in redshift and enables the discrimination between opposing viewpoints on HzRG formation and evolution - mature, red, passively evolving ellipticals or young, bursting irregulars - by permitting studies of the morphology of the stellar populations of HzRG environments. Work on imaging HzRGs in the NIR using the HST have been hampered by the failure of the Near-Infrared Camera and Multi-Object Spectrometer (NICMOS) in January 1999. From a sample of 19 $1.72$, but deeper imaging of a large sample is needed. The sample provided by natural guide-stars is insufficient and therefore the path will be provided by the next generation of AO systems, notably for Gemini and Keck, employing laser beacons." + }, + "0112/astro-ph0112154_arXiv.txt": { + "abstract": "Recent monitoring of Cyg X-1 with {\\em RXTE} revealed a period of intense flaring, which started in October of 2000 and lasted until March of 2001. The source exhibited some quite unusual behaviors during this period. The soft X-ray flux of the source went up and down three times on a timescale of about one month, as discovered by the ASM aboard RXTE, before finally returning to the normal level (of the hard state). The observed spectral and temporal X-ray properties of Cyg X-1 are mostly intermediate between the canonical hard and soft states. This is known previously for strong X-ray flares, however, we show that the source did enter a period that resembles, in many ways, a sustained soft state during the last of the three flares. We make detailed comparisons between this flare and the 1996 state transition, in terms of the observed X-ray properties, such as flux--hardness correlation, X-ray spectrum, and power density spectrum. We point out the similarities and differences, and discuss possible implications of the results on our understanding of the phenomena of flares and state transitions associated with Cyg X-1. ", + "introduction": "Cygnus X-1 has always been considered as an archetypical black hole candidate (BHC; reviews by Oda 1977, Liang \\& Nolan 1984, and Tanaka \\& Lewin 1995). Its observed spectral and temporal X--ray properties have, therefore, often been used to distinguish BHCs from their neutron star counterparts. Though flawed, this approach has resulted in the discovery of many BHCs whose candidacy is subsequently confirmed by dynamical measurements of the mass of the compact object. Cyg X-1 shows two distinct spectral states: the hard state and the soft state. The source spends most of the time in the hard state where the soft X-ray (2-10 keV) luminosity is relatively low and the power-law X-ray spectrum relatively flat. Once every few years, Cyg X-1 undergoes a transition from the hard state to the soft state. It remains in the soft state for weeks to months before returning to the hard state. During such a spectral state transition, the power-law spectrum ``pivots'' around 10--20 keV, causing a significant increase in the soft X-ray flux but a decrease in the hard X-ray ($\\gtrsim 20$ keV) (e.g., Zhang et al. 1997). The source has only occasionally been observed in the soft state (Oda 1977; Liang \\& Nolan 1984; Cui et al. 1997a) and is thus not as well studied as in the hard state. Besides the difference in the spectral properties of the source between the two states, the temporal properties are also quite different. The power density spectrum (PDS) of Cyg X-1 in the hard state can be characterized by a white noise component at low frequencies and a power-law component at high frequencies, with the characteristic ``break frequency'' in the range of $\\sim$0.04-0.4 Hz (reviews by van der Klis 1995 and Cui 1998). In the soft state, however, the PDS is dominated by ``$1/f$'' noise (Cui et al. 1997a). For both states, the PDS shows additional features at a few Hz (e.g., Belloni \\& Hasinger 1990; Cui et al. 1997a). Moreover, the X-ray flux of Cyg X-1 is known to vary on all timescales down to at least milliseconds; the X-ray variability seems stronger in the hard state than in the soft state. Cyg X-1 is also known to experience frequent X-ray flares. While most flares last for less than a day, some can last for weeks to months, during which additional short-duration flares may also occur. Studies have shown that at least some of the major flares are similar to state transitions (Pottschmidt et al. 2001; Feng, Ertmer, \\& Cui, in preparation). Therefore, the distinction between major flares and state transition can sometimes be quite ambiguous. The X-ray flares of Cyg X-1 are poorly understood due to the lack of high-quality data. In this {\\it Letter}, we present results from recent observations of Cyg~X-1 during a peculiar period of intense flaring activities with the {\\it Rossi X-ray Timing Explorer} (RXTE). ", + "conclusions": "We observed Cyg X-1 during an unusual period when three consecutive major X-ray flares occurred over a period of about five months. Following each flare, the source did seem to completely return to the hard state before the next one started, as indicated by the flux-hardness correlation. The correlation also indicates that Cyg X-1 entered a sustained period during the last flare that is very similar to the 1996 soft state. This is supported by striking similarities between the observed spectral and temporal properties of the source during that time period and those during the 1996 soft state (although there are some differences). Therefore, we argue that the physical process that is responsible for triggering a flare or a state transition is likely to be the same. The difference between the two types of phenomena seems mostly quantitative. For instance, the soft state lasted only for about a week during the flare, compared to more than two months in the 1996 episode This difference might be due to the difference in the change of physical quantities, such as the fraction of accretion energy dissipated in the disk. Although only an empirical model was used for spectral studies, the results do clearly reveal the evolution of the source, thanks to the improved coverages of the transitions. It was known previously that a broken power law would be required to model the spectrum of Cyg X-1 at high energies (Ebisawa et al. 1996; Cui et al. 1997a), as opposed to a simple power law for most BHCs (Tanaka \\& Lewin 1995). Not only did we confirm this in our investigation, we also quantified the evolution of this component throughout the transitions. Moreover, we confirmed the presence of a ``settling period'' after the low-to-high transition or before the high-to-low transition, as suggested by Cui et al. (1997a), when the soft flux is at the soft-state level but the spectral and timing properties are still intermediate between the hard and soft states. Such a period manifests itself prominently in the PDS shape. Finally, throughout the flare the observed 5--200 keV flux does not vary significantly. This is again consistent with previous results from a study of the 1996 state transition (Zhang et al. 1997). We conclude that we can perhaps learn a great deal about the origin of rare state transitions of Cyg X-1 from studying more frequently occurring X-ray flares. We have tried to establish a connection between the two types of phenomena in this investigation. The next step will be to carry out, in a more systematic manner, detailed comparisons between a range of X-ray flares and state transitions (Feng, Ertmer, \\& Cui, in preparation)." + }, + "0112/astro-ph0112044_arXiv.txt": { + "abstract": "We have compared simultaneous Ryle Telescope radio and Rossi X-Ray Timing Explorer X-ray observations of the galactic microquasar GRS~1915+105, using the classification of the X-ray behaviour in terms of three states as previously established. We find a strong (one--to--one) relation between radio oscillation events and series of spectrally hard states in the X-ray light curves, if the hard states are longer than $\\sim100$s and are ``well separated'' from each other. In all other cases the source either shows a low-level or a high level radio emission, but no radio oscillation events. During intervals when the source stays in the hard spectral state for periods of days to months, the radio behaviour is quite different; during some of these intervals a quasi-continuous jet is formed with an almost flat synchrotron spectrum extending to at least the near-infrared. Based on the similarities between the oscillation profiles at different wavelengths, we suggest a scenario which can explain most of the complex X-ray~:~radio behaviour of GRS~1915+105. We compare this behaviour with that of other black hole sources and challenge previous reports of a relation between spectrally soft X-ray states and the radio emission. ", + "introduction": "The black hole candidate GRS~1915+105 was the first Galactic source to display apparent superluminal motions of radio-emitting ejecta (Mirabel \\& Rodr\\'\\i guez 1994; Fender et al. 1999b; Rodr\\'\\i guez \\& Mirabel 1999). It is probably the best example to date of the strong coupling between the accretion disc and the jet in a black hole system (Pooley \\& Fender 1997; Eikenberry et al. 1998, 2000; Mirabel et al. 1998; Belloni, Migliari \\& Fender 2000). As an X-ray source GRS~1915+105 was discovered in 1992 with WATCH on board GRANAT (Castro-Tirado, Brandt \\& Lund 1992). Early observations, made with this instrument and with BATSE/CGRO, showed considerable X-ray variability. Since the launch of the Rossi X-ray Timing Explorer (\\emph{RXTE}) in 1995 the source has been continuously monitored with the All-Sky Monitor (ASM), and frequently observed with the Proportional Counter Array (PCA) and the High Energy X-ray Timing Experiment (HEXTE). From these observations the extraordinary X-ray variability displayed by GRS~1915+105 became clear (see for instance Greiner, Morgan \\& Remillard 1996; Belloni et al. 1997a,b; Muno, Morgan \\& Remillard 1999; Belloni et al. 2000a). Belloni et al. (2000a) found that the entire variability of the source is made up of a limited number of distinct variability classes; they identified twelve classes that each were observed to recur almost identically after intervals of months to even years. This behaviour is unique to GRS~1915+105, and is not observed in any other X-ray binary known to date. As in most black hole sources the X-ray spectrum of GRS~1915+105 can be modelled as a combination of two components: a soft (kT$\\sim$1 -- 2 keV) disc black body and a hard power law extending to $\\geq$100 keV, generally interpreted as representing the radiation from the accretion disc and a Comptonizing region (`corona'), respectively. Belloni et al. (2000a) found that the complex X-ray variability of GRS~1915+105 can be reduced to the transitions between three basic spectral states which they called A, B and C. The spectrally soft states A and B correspond to an observable inner accretion disc with different temperatures: in State B the inner temperature ($\\sim2.2$ keV) is higher compared to State A ($\\sim1.8$ keV). For the spectrally hard State C the inner part of the accretion disc is either missing or just unobservable, and the 2 -- 25 keV spectrum is dominated by a power law of spectral index $\\Gamma\\sim$1.3 -- 2.4. The description of the behaviour of GRS~1915+105 as transitions between these three spectral states differs from one based on the canonical black hole states (for a review on the canonical black hole states, see for instance Tanaka \\& Lewin 1995 and van der Klis 1995), as it is based on the positions in the X-ray colour-colour diagram, and not on the combined spectral and timing characteristics. There are, however, indications that the States A, B and C are related to the High State (HS), Very High state (VHS) and the Low State (LS) respectively (Belloni et al. 2000a and references therein). The description of the complex X-ray behaviour of GRS~1915+105 as the transitions between three basic spectral states, and the classification of the light curves into twelve different variability classes will be discussed in some more detail below. GRS~1915+105 is also amongst the most spectacular radio sources in our Galaxy, displaying apparent superluminal motions (Mirabel \\& Rodr\\'\\i guez 1994; Fender et al. 1999b; Rodr\\'\\i guez \\& Mirabel 1999) and radio oscillations which are clearly related to quasi-periodic dips in the X-ray light curve on time-scales of tens of minutes (Pooley \\& Fender 1997; Mirabel et al. 1998). The synchrotron emission of these oscillations extends well beyond the radio to the millimeter and near-infrared regimes (Fender et al. 1997; Eikenberry et al. 1998, 2000; Mirabel et al. 1998; Fender \\& Pooley 1998, 2000). The power required to generate these repeated synchrotron events is likely to be a very significant fraction of the entire accretion energy of the system (Fender \\& Pooley 2000, Meier 2001). There is evidence that the power into the jet (assuming that a jet is the origin of the radio oscillations) may be anticorrelated with the accretion rate as inferred from X-ray spectral fits (Belloni et al. 2000b). Belloni et al. (1997a,b) suggested that the X-ray dips correspond to the disappearance of the inner ($\\leq 100$ km) of the accretion disc. The relation of the radio oscillations to these dips led to the assertion that at least some of the inner disc which had `disappeared' had in fact been accelerated and ejected from the system, and had not simply fallen into the black hole. Based on the similarity between oscillation decays at different wavelengths, Fender et al. (1997) came to the conclusion that material is being ejected in the form of jets where the synchrotron emitting ejecta are predominantly subjected to adiabatic expansion losses. Whatever the interpretation, the coupling between radio--mm--infrared emission which is believed to arise in the jet, and the X-ray emission from the disc and corona, is very clear, and GRS~1915+105 is our most accessible route to date to an understanding of the connection between accretion and ejection around a black hole. In this paper we cross-correlate the entire database of radio observations with the Ryle Telescope (RT) at 15 GHz over the period 1996 -- 1999, with pointed X-ray observations with the \\emph{RXTE}. In addition we use some data from the Green Bank Interferometer (GBI) variable source monitoring program. The next two sections deal with the observation analysis and the X-ray~:~radio overlap. In Section~\\ref{sec:obs} we present our results by discussing a selection of representative observations in detail. In Sections~\\ref{sec:emp} and ~\\ref{sec:platstates} the empirical characteristics of the X-ray~:~radio correlation are discussed, and Section~\\ref{sec:delays} deals with the profiles of the radio oscillation events. In Section~\\ref{sec:discussion} we discus a possible scenario which explains the correlated X-ray~:~radio behaviour, and finally in Section~\\ref{sec:conclusions} we present our conclusions. ", + "conclusions": "\\label{sec:conclusions} \\begin{table} \\caption{The emipircal X-ray~:~radio correlation in GRS~1915+105.\\label{t4}} \\begin{tabular}{cc} \\hline Radio Behaviour\t&X-ray behaviour\\\\ \\hline \\hline Radio oscillation events &repeating, hard, well\\\\ &separated $\\ga100$s State C\\\\ \\hline Low level emission & adjacent, $\\la100$s \\\\ &State C or A, B\\\\ \\hline Plateau Jet &long uninterrupted\\\\ &State C\\\\ \\hline Pre-, Post-plateau flares &State transitions?\\\\ \\hline \\end{tabular} \\end{table} This work has clearly established that radio emission from GRS~1915+105 is intimately related to the presence of hard (power-law-dominated) intervals in the X-ray light curves. This in turn physically implies a clear relation between a radiatively inefficient flow close to the black hole, and a synchrotron-emitting outflow or jet. In Table~\\ref{t4} we summerize the empirical X-ray~:~radio correlation, moreover we find that: \\begin{itemize} \\item{The (high flux) radio oscillation events are closely related to a series of long (or low count rate), well-separated, spectrally hard State C intervals in the X-ray light curve.} \\item{The State C intervals must be at least 100s in length and have a recurrence time that allows for the individual radio oscillation events to be observed.} \\item{We confirm a one-to-one relation between radio oscillation events and series of long State C intervals: each State C interval produces a new ``radio flare''.} \\item{Short ($\\la100$s) State C intervals and State A and B intervals produce a similar low-level radio emission.} \\item{The radio oscillation profiles seem to be determined by particle injection/acceleration (rise phase) and adiabatic expansion losses (decay phase), which are both frequency independent effects.} \\item{During class $\\chi_1$ and $\\chi_3$ (plateau states) we find a quasi-continuous jet with a flat synchrotron spectrum (extending to at least the near-infrared) and a flux level of $\\sim100$ mJy, an anticorrelation between the radio flux and the soft (ASM) X-ray flux, and a correlation between the radio flux and the ASM HR$_{2}$ colour.} \\item{During class $\\chi_2$ we also find a flat synchrotron spectrum, but a flux level which is about a factor 10 lower, a correlation between the radio flux and the soft (ASM) X-ray flux and a anticorrelation between the radio flux and the ASM HR$_{2}$ colour.} \\end{itemize} Based on this we suggest a simple scenario which describes the complex X-ray~:~radio behaviour of GRS~1915+105: \\begin{itemize} \\item{During State C a more or less continuous ejection of relativistic particles takes place;} \\item{The length (or depth, as longer also means deeper) of the State C interval determines the strength of the radio emission: longer State Cs have a higher radio flux;} \\item{The separation between the State Cs determines the shape of the radio light curve: if the intervals are to close together one observes a ``stable'' radio flux level, while if the separation is large enough for the individual radio flares to be observed one finds a radio light curve dominated by oscillations.} \\end{itemize} Within the simple scenario we can explain: \\begin{itemize} \\item{the radio oscillation events: series of well separated, long State Cs;} \\item{the radio plateau jets: uninterrupted State Cs;} \\item{the low-level radio emission during the short State Cs: short, near-contiguous State Cs.} \\end{itemize} All these types of radio emission seem to be related to the same physical mechanism operating on different scales. However, this scenario does not explain: \\begin{itemize} \\item{The disproportional higher radio flux in the case of the long State C compared to the short State C intervals (although cooling by soft X-ray photons might be responsible);} \\item{The low-level radio emission in the case of State A and B, which cannot be explained based on argument of a continuous ejection during State C;} \\item{The radio plateau flares observed prior to and just after the plateau states, which seem to be related to the Soft -- Hard and Hard -- Soft X-ray state transitions as also observed in SXTs;} \\item{The lower radio emission during class $\\chi_2$ compared to classes $\\chi_1$ and $\\chi_3$.} \\end{itemize} With respect to the low-level radio emission found during States A, B and short State C intervals, we cannot rule out the possibility that the emission is a relic from a previous large radio event. The most important result is probably the establishment of the State C -- radio relation, and in particular the suggested one-to-one relation between the long State C intervals and the radio oscillation events. Because GRS~1915+105 has a much larger coverage in the radio than in the X-rays (PCA observations), this intimate relation can provide us with much more information about the behaviour of the source its behaviour than is revealed to us by the X-rays alone. However, due to the time delay between the radio and the X-rays, the disc -- jet connection is probably better studied in the X-ray and IR bands. The analysis presented here, together with previous reports on other black hole X-ray binary systems such as GX~339$-$4 (Fender et al. 1999a) and XTE~J1550-564 (Corbel et al. 2001), provides strong evidence that hard, power law-dominated states can support a radio-emitting outflow. This suggest, as discussed in Section~\\ref{sec:mhd}, that MHD effects could be responsible for the production and/or confinement of the jets found in these systems." + }, + "0112/astro-ph0112334_arXiv.txt": { + "abstract": "Optical broad-band polarimetry is presented for near-infrared color-selected active galactic nuclei (AGN) classified as quasi-stellar objects (QSOs) based on their $K_s\\/$-band luminosity. More than 10\\% of a sample of 70 QSOs discovered in the Two Micron All Sky Survey (2MASS) with $J - K_s > 2$ and $M_{K_s} < -23$ show high broad-band linear polarization ($P > 3$\\%), and values range to a maximum of $P \\sim 11$\\%. High polarization tends to be associated with the most luminous objects at $K_s\\/$, and with QSOs having the highest near-IR--to--optical flux ratios. The 2MASS QSO sample includes objects possessing a wide range of optical spectral types. High polarization is seen in two of 22 broad emission-line (Type~1) objects, but $\\sim 1/4$ of the QSOs of intermediate spectral type (Type~1.5--1.9) are highly polarized. None of the nine QSOs classified as Type~2 exhibit $P > 3$\\%. It is likely that the unavoidable inclusion of unpolarized starlight from the host galaxy within the observation aperture results in reduced polarization for the narrow emission-line objects. The high polarization of 2MASS-discovered QSOs supports the conclusion inferred from their near-IR and optical colors, that the nuclei of many of these objects are obscured to some degree by dust. Correlations between optical polarization and near-IR luminosity and color imply that the dominant polarizing mechanism in the sample is scattering of AGN light into our line of sight by material located in relatively unobscured regions near the nucleus. The broad-band polarization properties of the 2MASS QSO sample are compared to those of other, predominantly radio-quiet, QSOs and are found to be consistent with the idea that the orientation of AGN to the line of sight plays a major role in determining their observed properties. ", + "introduction": "Despite several decades of intense scrutiny, our understanding of active galactic nuclei (AGN) remains incomplete. Fundamental issues such as the number density of AGN, their contributions to the far-infrared and X-ray backgrounds, the role of orientation to our line of sight, the processes that form and power the nuclei, and their evolution still need to be resolved. Recent large-scale X-ray, optical, and near-infrared surveys have begun to reveal new populations of AGN with number densities comparable to, or exceeding, the number density of AGN derived from previous optical and radio surveys \\citep[e.g.,][]{sanders96}. In particular, the Two Micron All Sky Survey \\citep[2MASS;][]{skrutskie97} is finding large numbers of AGN that have been missed by traditional UV/optical search techniques. New survey wavebands can uncover past observational biases and reveal a wider range of AGN properties. Polarization is an important property that can be used to explore the physics and structure of AGN, and polarimetry yields valuable information about the various classes of active nuclei. For instance, early surveys of the most luminous known AGN -- quasi-stellar objects (QSOs) and quasi-stellar radio sources (quasars) -- showed that they could be divided into two classes by their optical polarization properties \\citep{moore84}. The vast majority of objects, both radio and optically selected, show low optical polarization ($P < 3$\\%; \\citealt{stockman84}). About 1\\% of the sample exhibited higher and variable polarization \\citep{moore81}, in part earning them the designation of optically violent variable quasars (OVVs). Almost all of these objects are associated with compact radio sources, and their variability behavior, both in flux and in polarization, clearly signifies that optically thin synchrotron radiation dominates the UV/optical continua. Identification of the polarization mechanism(s) in low-polarization QSOs has not been as straightforward largely because of the low levels of polarization encountered. Except for a handful of broad absorption-line QSOs (BALQSOs; \\citealt{moore84,schmidt99}) and QSOs discovered by the {\\it Infrared Astronomical Satellite\\/} ({\\sl IRAS\\/}) \\citep{hines94}, only the OVVs exhibit polarization above the traditional high-polarization threshold of $P = 3$\\%. A good illustration of the difficulty in identifying the polarizing mechanism in most QSOs is provided by the blue-excess Palomar-Green (PG) sample \\citep{schmidt83}. \\citet{berriman90} have shown that this predominantly radio-quiet sample has a mean intrinsic optical polarization of only 0.5\\% and the maximum observed polarization is 2.5\\%. Unambiguous statistical tests designed to determine the source of the polarized flux have been extremely hard to obtain, even though high quality broad-band polarimetry exists for the entire PG QSO sample. This is not the case for the the lower luminosity analogs to radio-quiet QSOs, the Seyfert~1 galaxies. These objects show overall polarization properties similar to the PG QSOs, but their polarization distribution extends to higher values. The high-polarization Seyfert~1 ``tail'' better facilitates comparison of the polarization distribution with dust extinction and dust emission indicators, and led \\citet{berriman89} to conclude that the primary polarizing mechanism in Seyfert~1s is the scattering of nuclear light by dust located within the narrow emission-line region (NLR). Variations of this mechanism have been successful in explaining the polarization of several Seyfert~2 galaxies and showing that these objects would look like Seyfert~1s if viewed from sightlines that intercept the scattering material and the NLR \\citep[e.g.,][]{antonucci85,millgood90,tran95b, wilkes95}. {\\sl IRAS\\/} found a small number of infrared-luminous AGN which suggest the existence of a population of red, optically-obscured radio-quiet QSOs \\citep{beichman86,low88,low89}. Spectropolarimetry of several of these ``hyperluminous'' infrared galaxies (HIGs) has confirmed that their tremendous IR luminosities are powered in large part by QSOs hidden from our direct view, and that polarization arises from scattering of nuclear light by material with a relatively unobstructed view of the AGN \\citep{wills92,hines95a,hines95b,hines99,hines01,goodrich96, young96,tran00}. A similar situation has been found for several radio galaxies \\citep{alighieri94,jannuzi95,cohen99}. For these objects, light scattered by dust located up to several kiloparsecs from the nucleus reveals the rest-frame ultraviolet spectrum of a QSO hidden from our direct view. \\citet{cutri01} have initiated a near-infrared, high galactic latitude survey for AGN using 2MASS. The survey is designed to find new red, radio-quiet objects and has, in fact, revealed a large population of previously unidentified AGN. The new sample generally consists of low-redshift ($z < 0.7$) objects with redder near-IR colors (a criterion of the object selection) and much higher near-IR--to--optical flux ratios than optically selected AGN. Many of these AGN can be classified as QSOs in their own right since their $K_s\\/$-band luminosities are comparable to those of optically selected QSOs. Higher levels of optical polarization relative to UV-excess QSOs are expected from the 2MASS sample if, as with the IR-selected HIGs, many of these objects are obscured along our line of sight. Nuclear light scattered by material (dust and/or electrons) distributed around the AGN will make a larger contribution to the total observed flux than if, as inferred for optically selected QSOs, we have a direct, unobstructed view to the bright unpolarized nucleus. Given an asymmetric distribution and/or illumination of the scattering material, high polarization can result. The 2MASS AGN survey presents a large, well-defined sample of QSOs that invites comparisons with AGN selected by other means. In this paper we present observations of 2MASS QSOs aimed at determining their broad-band optical polarization properties. We test the conclusion reached by \\citet{cutri01}, based on the near-IR colors of the sample, that some of the objects in the general low-$z\\/$ QSO population found by 2MASS are obscured along our line of sight. In addition, the polarization distribution of the new near-IR--selected QSOs is compared to three largely radio-quiet QSO samples (PG, BALQSO, and {\\sl IRAS\\/} HIGs) for which polarization data are available. ", + "conclusions": "The broad-band optical polarization of QSOs selected by their near-IR colors are consistent with the view that many of these objects are obscured to some degree by dust. Over 10\\% of a sample of 70 ``red'' QSOs discovered by 2MASS \\citep{cutri01} shows high polarization ($P > 3$\\%), and values range up to $\\sim$11\\%. These high levels of polarization are not present among objects in standard UV-excess catalogs of QSOs except for OVVs and a small number of BALQSOs. The polarization of 2MASS QSOs is correlated with near-IR luminosity and color, and with near-IR--to--optical flux ratio in the sense that more luminous and redder objects tend to show high polarization. These trends suggest that the polarization arises from the scattering of nuclear light by material located close to a partially obscured AGN. Spectropolarimetry of two highly polarized 2MASS QSOs confirms that scattering is the dominant polarizing mechanism in these objects with the scattering material located close to, or within, the NLR \\citep{smith00a}. Determining the nature of the scattering material (either dust and/or electrons), however, likely awaits careful modeling of high-quality spectropolarimetry at UV, optical, and infrared wavelengths \\citep[see e.g.,][]{hines01}. The near-IR color selection of the 2MASS AGN survey is unbiased with regard to optical spectral type, and we find that high polarization is found among nearly all types, especially those classified as Type~1.5--1.9 by \\citet{cutri01}. Like optically selected QSOs, few of the broad emission line-dominated (Type~1) QSOs have $P > 3$\\%. This is presumably because the obscuration of direct AGN light is sufficiently low that the relative contribution of any scattered component is small. High polarization is not observed in Type~2 QSOs in this sample, although their colors and $K_s\\/$ luminosities are consistent with these objects being among the most heavily obscured in the sample. Dilution of scattered, polarized light by starlight from the host galaxy is suggested as the reason for low observed optical polarization in this case. This assertion is consistent with the fact that most Seyfert~2 galaxies that are highly polarized in the UV show $P < 3$\\% for $\\lambda > 5000$~\\AA . Ultraviolet polarimetry and high quality imaging provide obvious tests of the effect of host galaxy starlight on the polarization of 2MASS QSOs. The distribution of optical polarization for the 2MASS QSOs further substantiates the trend that higher polarization is seen in QSO samples that have higher apparent IR--to--optical luminosity ratios ($L_{\\rm IR}\\//L_{\\rm opt}\\/$). The 2MASS objects are not as extreme in $L_{\\rm IR}\\//L_{\\rm opt}\\/$ as the {\\sl IRAS\\/} HIGs, and neither is their general level of polarization. On the other hand, the 2MASS sample is much more highly polarized than the PG QSOs, which generally have the lowest $L_{\\rm IR}\\//L_{\\rm opt}\\/$. This is presumably because our view of the nuclear region is not hindered by dust in UV-excess AGN. Although the 2MASS and BALQSO distributions in $P\\/$ are fairly similar, corrections for starlight and for differing observing bandpasses are likely to better differentiate these two samples and thereby increase the significance of the correlation between $P\\/$ and $L_{\\rm IR}\\//L_{\\rm opt}\\/$. This correlation implies that orientation effects and/or the dust covering factor are largely responsible for differences in the optical and IR properties between various samples of high-luminosity AGN. Alternatively, the near-IR--selected QSOs may represent a different evolutionary stage than the optically selected, UV-excess QSOs. Further orientation-independent investigations of these samples, such as imaging of host galaxies and environments, extended radio emission, and IR observations using {\\sl SIRTF\\/} will be necessary to distinguish between evolutionary and orientation effects. Finally, the polarization properties of the 2MASS QSOs suggests that this new near-IR--selected sample may include a large proportion of low-redshift BALQSOs. Ultraviolet observations can directly test this possibility, and together with X-ray measurements, investigate the nature of the obscuring material." + }, + "0112/astro-ph0112428_arXiv.txt": { + "abstract": "{High-resolution IRAS maps are used to search for the presence of stellar-wind bow-shocks around high-mass X-ray binaries (HMXBs). Their high space velocities, recently confirmed with {\\it Hipparcos} observations, combined with their strong stellar winds should result in the formation of wind bow-shocks. Except for the already known bow-shock around Vela~X-1 (Kaper et al. \\cite{kaper97}), we do not find convincing evidence for a bow-shock around any of the other HMXBs. Also in the case of (supposedly single) OB-runaway stars, only a minority appears to be associated with a bow-shock (Van Buren et al. \\cite{vanburen95}).\\\\ We investigate why wind bow-shocks are not detected for the majority of these OB-runaway systems: is this due to the IRAS sensitivity, the system's space velocity, the stellar-wind properties, or the height above the galactic plane? It turns out that none of these suggested causes can explain the low detection rate ($\\sim$40~\\%). We propose that the conditions of the interstellar medium mainly determine whether a wind bow-shock is formed or not. In hot, tenuous media (like inside galactic superbubbles) the sound speed is high ($\\sim$100~\\kms), such that many runaways move at subsonic velocity through a low-density medium, thus preventing the formation of an observable bow-shock. Superbubbles are expected (and observed) around OB associations, where the OB-runaway stars were once born. Turning the argument around, we use the absence (or presence) of wind bow-shocks around OB runaways to probe the physical conditions of the interstellar medium in the solar neighbourhood. ", + "introduction": "OB-runaway stars are massive, OB-type stars with high space velocities. This small ($\\sim$20~\\% of the O stars, less than 5~\\% of the early B stars), but significant excess of high-velocity objects suggests a systematical production of fast movers (Feast \\& Shuttleworth \\cite{feast}, Stone \\cite{stone}). Among the OB runaways, velocities well above 100~\\kms\\ have been observed (Gies \\& Bolton \\cite{gies86}, Conlon et al. \\cite{conlon}, Blaauw \\cite{blaauw93}), which is roughly ten times the average space velocity of ``normal'' OB stars in the Milky Way. Due to their high velocities, OB-runaways may have travelled great distances since their formation. The criteria used to classify OB stars as ``runaways'' are (Blaauw \\cite{blaauw61}): (i) A high space velocity; often a threshold space velocity of 30~\\kms\\ is used, roughly corresponding to three times the typical space velocity of OB-type stars. (ii) The reconstructed path of the runaway should start in a ``parent'' OB association. Since the definition by Blaauw, the term {\\it OB runaway} has often been used to include all of the high-velocity OB stars, irrespective of the identification of a parent OB-association. A large distance above the galactic plane can also be used as an indication for the runaway-nature of a massive star (e.g. Van Oijen \\cite{vanoijen}). The two most popular scenarios for the production of runaways seem to operate at roughly the same rate (Hoogerwerf et al. \\cite{hoogerwerf}) and are both capable of producing runaways with velocities up to 200~\\kms\\ (Portegies-Zwart \\cite{portegies}, Leonard \\cite{leonard}): (i) The {\\it binary supernova scenario} (Blaauw \\cite{blaauw61}, Zwicky \\cite{zwicky}), where a massive star in a binary receives a high space velocity after the supernova explosion of its (initially) more massive companion; (ii) The {\\it cluster ejection scenario} (Poveda et al. \\cite{poveda}), where dynamical interactions in a compact cluster result in the ejection of one or more of its members. Especially encounters between binary systems are effective in producing high-velocity objects (e.g. Mikkola \\cite{mikkola}). The current version of the binary supernova scenario includes a phase of mass transfer (Van den Heuvel \\& Heise \\cite{VdhH72}): when the initially most massive star fills its Roche lobe (e.g. when becoming a supergiant) mass is transfered to its companion which eventually will become the most massive star in the system. As a consequence, the system has a high probability to remain bound after the supernova explosion of the initially most massive star (Boersma \\cite{boersma61}). If it remains bound, the binary, now consisting of a massive, rejuvenated main-sequence star and a compact remnant (a neutron star or a black hole), will travel through space with a high velocity. As soon as the massive star becomes a supergiant, material from its dense stellar wind (or through Roche-lobe overflow) is intercepted by the gravitational field of the compact star and accretes, powering a strong X-ray source: a high-mass X-ray binary (HMXB). This scenario thus predicts that all HMXBs are runaway objects. Recently, this prediction has been confirmed with {\\it Hipparcos} observations (Chevalier \\& Ilovaisky \\cite{chevalier}, Kaper et al. \\cite{kaper99}). This also suggests that several supposedly single OB runaways might have a, so far undetected, compact companion. For a review on HMXBs and OB-runaway stars, see Kaper (\\cite{kaper01}). \\begin{table*}% \\centering \\caption{Prediction of the angular separation between wind bow-shock and HMXB. The HMXB peculiar velocity (i.e. corrected for differential galactic rotation and peculiar solar motion) in the radial and tangential direction is listed in column 6 and 7, respectively; the resulting space velocity $v_{\\rm space}$ is given in column 8. The first five columns list the identification of the X-ray source, the name and spectral type of the OB companion (if known), the distance $r$, and the height $z$ above the galactic plane. The predicted (and observed, final column) angular distance between OB star and wind bow-shock are given in the last columns. For the HMXBs that have not been observed with {\\it Hipparcos} $(v_{\\rm tan})_{pec}=$ 50~\\kms\\ has been adopted. If a question mark is added, infrared emission is detected, but no clear arclike structure could be resolved. The prediction of the stand-off distance is based on stellar-wind parameters taken from Puls et al. (\\cite{puls}) and Kaper et al. (\\cite{kaper98}); for Be stars it was assumed that $\\dot{\\rm M} = 10^{-8} {\\rm M}_{\\odot}~{\\rm yr}^{-1}$ and $v_{\\infty}$ = 250~\\kms.} \\renewcommand{\\arraystretch}{1.4} \\setlength\\tabcolsep{5pt} \\begin{tabular}{lllcrrcc|cc|cc} \\hline\\noalign{\\smallskip} Object & & Spectral Type& $r$ & $z$ & $(v_{\\rm rad})_{pec}$ & $(v_{\\rm tan})_{pec}$ & $v_{\\rm space}$ & $d_{\\rm shock}^{\\rm pred}$ & $\\delta_{\\rm shock}^{\\rm pred}$ & bow-shock? & $\\delta_{\\rm shock}$ \\\\ & & & (kpc) & (pc) & (\\kms) &(\\kms) &(\\kms) &(pc) & $(')$ & & $(')$ \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 0114+650 & & B0.5 Ib & 3.8 & 171 & -9.6 & 29.4 & 31.0 & 2.6 & 2.4 & no & \\\\ 1223-624 & GX301-2 & B1.5 Ia & 5.0 & -3.0 & -1.6 & {\\em 50} & {\\em 50}& 2.8 & 1.9 & ? & 1.5(?) \\\\ 1700-377 & HD 153919 & O6.5 Iaf & 1.7 & 65 & -42.0 & 57.3 & 71.0 & 3.1 & 6.3 & ? & 8.0(?) \\\\ 1907+097 & & B0.5 Ib & 4.0 & 37 & & {\\em 50} &{\\em 50}& 0.6 & 0.5 & ? & 0.5(?) \\\\ 1538-522 & QV Nor & B0 Iab & 5.5 & 206 & -81.9 & {\\em 50} &{\\em 96} & 1.9 & 1.2 & no & \\\\ 1119-603 & Cen X-3 & O6.5 II-III & 8.0 & 47 & 16.3 & {\\em 50} &{\\em 53}& 1.7 & 0.7 & no & \\\\ 1956+350 & Cyg X-1 & O9.5 Iab & 2.5 & 135 & -10.9 & 41.4 & 42.8 & 5.0 & 6.9 & no & \\\\ 0900-403 & Vela X-1 & B0.5 Ib & 1.8 & 123 & -22.3 & 38.5 & 44.5 & 1.4 & 2.7 & yes & 1.0 \\\\ 0236+610 & V615 Cas & B0e & 2.0 & 42 & -30.2 & 22.5 & 37.7 & .14 & .24 & no & \\\\ 0535+262 & V725 Tau & O9.7e II & 2.0 & -90 & -41.5 & 57.5 & 70.9 & .07 & .12 & no & \\\\ 0352+309 & X Per & O9e III-IV & 0.8 & -235 & -32.0 & 13.4 & 34.7 & .51 & 2.2 & no & \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\end{tabular} \\label{tablehmxb} \\end{table*} The high space velocity and powerful stellar wind of the OB-runaway star will give rise to a strong interaction with the ambient medium. In the case of supersonic movement, a wind bow-shock is formed (Baranov et al. \\cite{baranov}). The wind bow-shock accumulates gas and dust from the interstellar medium which is heated by the ultraviolet radiation field of the OB star. The heated dust subsequently radiates at infrared wavelengths (Draine \\& Lee \\cite{draine}) and the shocked gas becomes visible in strong optical emission lines like H$\\alpha$ and [O~{\\sc iii}] 4959,5007~\\AA\\ (e.g.\\ Kaper et al. \\cite{kaper97}). Data obtained with the {\\it InfraRed Astronomical Satellite} (IRAS) has been used to search for the presence of wind bow-shocks around OB runaways by Van Buren \\& McCray (\\cite{vanburen88}) and Van Buren et al. (\\cite{vanburen95}). They detected excess infrared (60~$\\mu$m) emission for about 30~\\% of the 188 OB stars in their sample. For a significant fraction prominent arc-like features could be resolved using high-resolution maps (Noriega-Crespo et al. \\cite{noriega}). When spatially resolved, the shape of the bow-shock indicates the direction of motion of the system. The stand-off distance of the shock provides a constraint on the space velocity, the stellar-wind parameters, and the density of the ambient medium. In this study we investigate whether wind bow-shocks are as common for HMXBs as for (single) OB-runaway stars, by searching for their infrared emission in high-resolution maps produced from data obtained with the IRAS satellite. We find that also in the case of HMXBs only a minority form a wind bow-shock. We investigate whether the detection of a bow-shock relates to (i) the kinematical properties of the runaways; (ii) the height above the galactic plane; or (iii) the spectral type of the OB runaway. The remaining option, namely that the physical conditions of the ambient medium mainly determine whether a wind bow-shock is formed, seems to be the most likely one. If true, this provides the opportunity to use OB runaways as probes of the physical conditions of the interstellar medium. ", + "conclusions": "Based on our study of OB runaways and the phenomenon of wind-bow-shock formation we conclude that about 40~\\% of the OB-runaway stars produce wind bow-shocks. This fraction is a little bit higher than proposed by Van~Buren and coworkers, but takes into account the estimated reduction in IRAS sensitivity required to detect a bow-shock. We investigate why the majority of OB runaways do not produce a bow-shock. The spectral type (temperature, luminosity, mass-loss rate), space velocity and separation from the galactic plane seem to be of minor importance. We propose that the temperature and density of the ambient medium are the dominant factors. In hot bubbles the speed of sound is larger than the typical velocity of a runaway star ($\\sim$50~\\kms), so that runaways in these regions do not move supersonically and would not produce a bow-shock. OB-runaways produced through the binary supernova scenario should on average have a shorter kinematical age than those produced through the cluster ejection scenario, and thus have a higher probability to be still inside a hot bubble. For some runaways, e.g. Cyg~X-1, there is observational evidence that the system indeed is contained in a (super)bubble." + }, + "0112/astro-ph0112102_arXiv.txt": { + "abstract": "{ We investigate the error implied by the use of the Zel'dovich approximation to set up the initial conditions at a finite redshift $\\zi$ in numerical simulations. Using a steepest-descent method developed in a previous work (\\cite{paper2}) we derive the probability distribution $\\cP(\\dR)$ of the density contrast in the quasi-linear regime. This also provides its dependence on the redshift $\\zi$ at which the simulation is started. Thus, we find that the discrepancy with the exact pdf (defined by the limit $\\zi \\rightarrow \\infty$) is negligible after the scale factor has grown by a factor $a/\\ai \\ga 5$, for scales which were initially within the linear regime with $\\sigma_{\\rm i} \\la 0.1$. This shows that the use of the Zel'dovich approximation to implement the initial conditions is sufficient for practical purposes since these are not very severe constraints. ", + "introduction": "In usual cosmological scenarios, large-scale structures in the universe form through the growth of small initial density fluctuations by gravitational instability (e.g., \\cite{Peebles1}). Besides, in most cases of cosmological interest the amplitude of these perturbations increases at small scales, as in the CDM model (\\cite{Peebles3}). This leads to a hierarchical scenario of structure formation where smaller scales become non-linear first. They build small virialized objects which later become part of increasingly large structures. Thus, this halos give rise to galaxies or clusters of galaxies (depending on cooling processes). Unfortunately, this non-linear regime is very difficult to handle analytically so that N-body simulations are a key tool to understand the formation of large-scale structures. They are even more important when one tries to follow the evolution of baryons which involves many processes (star formation, cooling, radiative ionization,..). Therefore, it is important to get a good estimate of the accuracy of such numerical simulations. In this respect, an obvious source of error is the generation of initial conditions. Indeed, numerical simulations are initialized at a finite redshift $\\zi$ while they should be started at time $\\ti=0$ when the relevant scales are exactly ``linear''. In practice, one uses the Zel'dovich approximation (\\cite{Zel1}) to set up the initial conditions at $\\zi$. This correctly matches the exact density and velocity fields at linear order but the higher-order terms are not exact (as compared with the fields defined by the same linear growing mode initialized at $t=0$). This leads to a small error which can be made negligible by starting the simulations at a sufficiently large redshift $\\zi$, when the scales of interest are far within the linear regime (e.g., \\cite{Jus1}, \\cite{Bau1}). These transients induced by the Zel'dovich approximation were investigated in \\cite{Scoc1} for the first few order moments ($q \\leq 8$) of the density and velocity fields using a perturbative approach. However, this approach involves lengthy calculations which worsen at higher orders and it is not obvious how to get an estimate of the error for the probability distribution function (pdf) $\\cP(\\dR)$ of the density contrast from these moments. Therefore, in this article we show how one can apply to this problem a steepest-descent method developed in a previous work (\\cite{paper2}) which provides a rigorous derivation of $\\cP(\\dR)$ in the quasi-linear limit. This approach is quite general since it applies to Gaussian initial conditions (\\cite{paper2}) as well as to non-Gaussian initial conditions (e.g., \\cite{paper3}) or to the tails of $\\cP(\\dR)$ in both linear and non-linear regimes (\\cite{paper4}). Also, it is actually quite intuitive. Here we describe how it can be used to derive the dependence on the initial redshift $\\zi$ of the error induced on the pdf $\\cP(\\dR)$. It also yields the moments of the density field at any order. This article is organized as follows. First, in Sect.\\ref{Steepest-descent method} we recall the path-integral formulation which allows us to derive the pdf $\\cP(\\dR)$ in terms of initial conditions. Then, in Sect.\\ref{Zel'dovich initial conditions} we apply this approach to the gravitational dynamics defined by Zel'dovich initial conditions at a finite redshift $\\zi$. Finally, in Sect.\\ref{Numerical results} we present our numerical results. ", + "conclusions": "Thus, in this article we have shown how to derive the pdf $\\cP(\\dR)$ of the density contrast implied by the use of the Zel'dovich approximation to set up the initial conditions at a finite redshift $\\zi$. This allows us to obtain a rigorous quantitative estimate of the error due to this effect for the pdf $\\cP(\\dR)$ and the moments of the density contrast. Our results apply to large scales which are still in the quasi-linear regime. This is quite sufficient for practical purposes since it is clear that it makes no sense to use the Zel'dovich approximation at time $\\ti$ for scales which are already non-linear. Then, we found that the error for the pdf $\\cP(\\dR)$ associated with the Zel'dovich approximation at $\\zi$ is negligible after the scale factor has grown by a factor $a/\\ai \\ga 5$, for scales which were initially within the linear regime with $\\sigma_{\\rm i} \\la 0.1$. In low-density universes one needs a larger expansion factor ratio $a/\\ai$. In fact, to a good approximation the relevant quantity is the growth factor ratio $\\Dp(t)/\\Dp(\\ti)$, which should be larger than $5$. Thus, in order to use numerical simulations it is important to check that these constraints are indeed satisfied over the scales one is interested in. However, these conditions are not very severe, hence it appears that this effect is not the main source of error. Therefore, it is sufficient to use the Zel'dovich approximation to set up initial conditions and one does not need to use higher-order approximations." + }, + "0112/hep-ph0112351_arXiv.txt": { + "abstract": "We demonstrate by numerical flux calculations that neutrino beams producing the observed highest energy cosmic rays by weak interactions with the relic neutrino background require a non-uniform distribution of sources. Such sources have to accelerate protons at least up to $10^{23}$ eV, have to be opaque to their primary protons, and should emit the secondary photons unavoidably produced together with the neutrinos only in the sub-MeV region to avoid conflict with the diffuse $\\gamma-$ray background measured by the EGRET experiment. Even if such a source class exists, the resulting large uncertainties in the parameters involved in this scenario does currently not allow to extract any meaningful information on absolute neutrino masses. ", + "introduction": "In acceleration scenarios ultra high energy cosmic rays (UHECRs) with energies above $10^{18}\\,$eV are assumed to be protons accelerated in powerful astrophysical sources. During their propagation, for energies above $\\gtrsim50$ EeV ($1EeV = 10^{18} eV$) they lose energy by pion production and pair production (protons only) on the cosmic microwave background (CMB). For sources further away than a few dozen Mpc this would predict a break in the cosmic ray flux known as Greisen-Zatsepin-Kuzmin (GZK) cutoff~\\cite{gzk}, around $50\\,$EeV. This break has not been observed by experiments such as Fly's Eye~\\cite{Fly}, Haverah Park~\\cite{Haverah}, Yakutsk~\\cite{Yakutsk} and AGASA~\\cite{AGASA}, which instead show an extension beyond the expected GZK cutoff and events above $100\\,$EeV. However the new experiment HiRes~\\cite{Hires} currently seems to see a cutoff in the monocular data~\\cite{Hires2001}. Taking into account that all old experiments except perhaps AGASA do not have sufficient statistics in the highest energy region to settle the question, the existence of a possible cutoff remains unclear at the moment. The apparent absence of a cutoff especially in the AGASA data has in recent years triggered many theoretical explanations ranging from conventional acceleration in astrophysical sources to models invoking new physics such as the top-down scenarios in which energetic particles are produced in the decay of massive relics from the early Universe~\\cite{bs}. This enigma has also fostered the development of large new detectors of ultra-high energy cosmic rays which will increase very significantly the statistics at the highest energies~\\cite{reviews}. In bottom-up scenarios of UHECR origin, in which protons are accelerated in powerful astrophysical objects such as hot spots of radio galaxies and active galactic nuclei~\\cite{biermann}, one would expect to see the source in the direction of arrival of UHECRs, but above the GZK cutoff in general no suitable candidates have been found within the typical energy loss distance of a few tens of Mpc for the known electromagnetically or strongly interacting particles~\\cite{ssb,ES95}. Even assuming significant deflection by large scale extragalactic magnetic fields requires at least several sources~\\cite{ils} whose locations have not been identified yet. Moreover, recent observations of small scale clustering by the AGASA experiment~\\cite{AG1} suggest that sources of UHECR are point-like~\\cite{anis1,anis2}. This fact together with the lack of nearby sources favors the possibility of sources much further away than 100 Mpc, at redshifts of order unity. An additional motivation for this possibility comes from recently reported possible correlations of the arrival directions of observed UHECR above $\\simeq50\\,$EeV with certain classes of sources such as compact radio galaxies~\\cite{corr_radio} or BL Lacertae objects~\\cite{corr_bllac}. In the latter case it is still possible that the sources are located at moderate distances $z\\simeq0.1$. In this case photons with extremely high energies $E>10^{23}$ eV can propagate several hundred Mpc (constantly loosing energy) and can create secondary photons inside the GZK volume~\\cite{photons}. However, this model requires both extreme energies of primary photons and extremely small extra galactic magnetic fields (EGMFs) $B\\lesssim10^{-12}$~G. Moreover, if a correlation with any source at redshift $z > 0.2$ is found, this model will be ruled out. If sources of the highest energy cosmic rays are indeed at cosmological distances $z \\sim 1$, the only known mechanism not involving new physics except for neutrino masses assumes neutrinos as messenger particles: Charged particles accelerated in such sources give rise to a secondary neutrino beam which can propagate essentially unattenuated. If this neutrino beam is sufficiently strong it can produce the observed UHECRs within 100 Mpc by electroweak (EW) interactions with the relic neutrino background~\\cite{zburst1}. Specifically, if the relic neutrinos have a mass $m_\\nu$, Z-bosons, whose decay products can contribute to the UHECR flux, can be resonantly produced by ultra high energy (UHE) neutrinos of energy $E_\\nu\\simeq M_Z^2/(2m_\\nu)\\simeq4.2\\times10^{21}\\,{\\rm eV}\\,({\\rm eV}/m_\\nu)$. However, this ``Z-burst'' mechanism is severely constrained by at least two types of observational data: First, there are upper limits on the UHE neutrino flux, based on the non-observation of horizontal air showers by the old Fly's Eye experiment~\\cite{baltrusaitis} or by the AGASA experiment \\cite{AgAsA} and from the non-observation of radio pulses that would be emitted from the showers initiated by the UHE neutrinos on the moons rim~\\cite{goldstone}. Second, even if the sources exclusively emit neutrinos, the EW interactions also produce photons and electrons which initiate an electromagnetic (EM) cascade which transfers the injected energy down to below the pair production threshold for photons on the CMB~\\cite{bs}. The cascade thus gives rise to a diffuse photon flux in the GeV range which is constrained by the flux observed by the EGRET instrument on board the Compton $\\gamma-$ray observatory~\\cite{egret}. Reproducing the observed UHECR flux by the Z-burst mechanism under these two constraints has been shown to in general require local relic neutrino over-densities in order to increase the local UHECR flux. These over-densities turn out to be much higher than values 2--3 which would be expected from the over-density in the local supercluster~\\cite{zburst2}. In order to avoid this difficulty one can suppose that the Z-burst mechanism is responsible only for part of the UHECR flux~\\cite{fkr}. In this case, one can reduce both primary neutrino and secondary photon fluxes and obey all existing limits. However, the price for this is to explain only a part of the UHECR events by the Z-burst mechanism and the necessity for a second source mechanism for UHECRs. Furthermore, Ref.~\\cite{fkr} claims that already the present data provides possible evidence for the relic neutrino background and starts to constrain the absolute neutrino mass, a possibility that has recently been discussed in principle in Ref.~\\cite{pw}. This claim is based on tuning many unknown parameters such as the value of the EGMF, the universal radio background (URB) which governs pair production of UHE $\\gamma-$rays, and the neutrino source distribution. Also, Ref.~\\cite{fkr} did not take into account propagation of UHE photons, instead assuming that all photons are down-scattered into the GeV region. In addition, simply due to the much larger statistics at lower energies, the quality of the fits performed in Ref.~\\cite{fkr} is dominated by the low-energy background component. Finally, Ref.~\\cite{fkr} assumed that the sources do not emit any $\\gamma-$rays, although the $\\gamma-$ray energy fluence produced by pion production of accelerated nuclei should be comparable to the produced neutrino fluence, as will be discussed in Sec.~\\ref{phot+nu}. In the present paper we show that for all neutrino masses in the range $ 0.07 {\\rm eV}\\leq m_\\nu\\leq 1 {\\rm eV}$ one can find parameters that fit the UHECR observations with comparable quality. We therefore conclude that, at least at the current state of knowledge, it is impossible to extract evidence for the relic neutrino background or even best fit values for absolute neutrino masses from UHECR data. We do not consider in the present paper neutrino interaction channels with multiple $W^{\\pm}$ and/or $Z^0$ production. These channels could be important in case of neutrino masses $m \\gtrsim 3$ eV \\cite{Fargion}, which however are strongly disfavored by considerations on large scale structure formation \\cite{Fukugita}. By detailed numerical flux calculations we show that a non-uniform source distribution allows the Z-burst mechanism to explain the UHECR flux without substantial relic neutrino over-densities. However, this only works if the sources exclusively emit neutrinos. Because isospin symmetry requires the energy fluence of neutrinos and $\\gamma-$rays produced by hadronic charged primary interactions in the source are comparable, this will require the photons to be down-scattered below the GeV range within the source. ", + "conclusions": "The Z-burst mechanism where the highest energy cosmic rays are produced by neutrino beams interacting with the relic neutrino background only works with sources exclusively emitting neutrinos in the ultra-high energy regime. In order to avoid conflict with the known diffuse backgrounds of $\\gamma-$rays, these sources should emit photons only in the sub-MeV region. In addition, they should trap primary protons in order to avoid an excessive nucleon flux from the source, and should be able to accelerate these protons up to $E^p_{\\rm max} \\gtrsim 10^{23} \\,{\\rm eV}\\,({\\rm eV}/m_\\nu)$. None of the astrophysical acceleration models existing in the literature seems to meet this requirement. Under the assumption that such an extreme source class nevertheless exists we have shown that the Z-burst mechanism can work without unrealistically high local relic neutrino over-densities if the neutrino sources are typically more abundant at present than in the past. Especially neutrino masses $m_\\nu\\lesssim0.5\\,$eV require a non-uniform source distribution $\\propto(1+z)^m$ with negative evolution factor, $m<0$, as is the case with BL Lacertae objects. The contribution to the UHECR flux from such a speculative extragalactic neutrino source class due to the Z-burst mechanism would exhibit a GZK-cutoff for nucleons and would be dominated by $\\gamma-$rays at higher energies. Furthermore, the required UHE neutrino fluxes are close to existing upper limits and should be easily detectable by future experiments such as Auger~\\cite{auger}, Euso~\\cite{euso}, RICE~\\cite{rice}, or by other radio detection techniques~\\cite{radhep}. The space of parameters characterizing neutrino sources and their evolution is highly degenerate when fluxes are fit to the observed UHECR fluxes. Since evidence of relic neutrinos and extraction of absolute neutrino masses requires conservative assumptions about these unknown parameters, we conclude that the current state of knowledge does not allow to extract any meaningful information on neutrino masses from UHECR data." + }, + "0112/astro-ph0112387_arXiv.txt": { + "abstract": "{\\ni We propose a theoretical model to explain the spectrum of the quasar PG1211+143 emitted in the Optical/X-ray bands. In particular, we suggest that the inner accretion disk may develop a warm, optically thick skin, which produces a profound emission feature observed in the soft X-ray band. This is well modelled with the Comptonized black body emission. The same warm, mildly ionized medium may also be responsible for the hard X-ray reflection and the presence of the iron $K_{\\alpha}$ line. However, in our model it still remains an open question, whether the seed photons for Comptonization come from the cold accretion disc or from the hotter plasma. High resolution spectroscopy available through the Chandra and XMM data may provide now an independent test of the physical conditions in the Comptonizing and reflecting warm skin. } \\bsk ", + "introduction": "The Narrow Line quasar PG1211+143 is a very good candidate to study the broad band emission and variability properties due to a rich sample of observational data available for this object. The source exhibits particularly strong hard optical spectrum and profound steep soft X--ray emission (Elvis, Wilkes, \\& Tananbaum 1985). It has been frequently argued that the optical/UV and soft X--ray emission form a single Big Blue Bump component which extends across the unobserved XUV band (Bechtold et al. 1987), and strongly dominates the bolometric luminosity. However, the nature of the soft X--ray emission below $\\sim 1$ keV is still not determined unambiguously, because any interpretation relies on extrapolating the spectrum over a decade in frequency through the unobserved XUV band. Detailed analysis of the broad band spectra of Seyfert 1 galaxy NGC~5548 (Magdziarz et al. 1998) led to the conclusion that the optical/UV spectrum is well modeled by an accretion disk, and the hard X--rays are reproduced by standard thermal Comptonization. However, an additional component is needed to model the soft X--ray source: the most viable candidate here is another, optically thick Comptonizing medium possibly associated with the transition between the disk and the hot plasma. Here we investigate if a similar model also applies to PG1211+143, and if this model is unique. We have analyzed the data form the ROSAT setellite, which span the period 1991-1993, ASCA observation form June 1993 and RXTE observation form August 1997. The ROSAT and ASCA spectral data were reduced using the standard software package, where for ASCA, we used data from all four detectors (SIS0/1 and GIS2/3) fitted simultaneously. For RXTE observation we use the spectrum derived from the summed Proportional Counter Array (PCA) data analyzed in a standard manner (including the background subtraction), using the {\\sc ftool} script {\\sc rex}. All analysis of the spectral X--ray data was done using the XSPEC software package, version 10.0 (Arnaud 1996). The details are given in Janiuk et al. 2001. \\bsk ", + "conclusions": "" + }, + "0112/astro-ph0112178_arXiv.txt": { + "abstract": "{We calculate the accretion disc temperature profiles, disc luminosities and boundary layer luminosities for rapidly rotating neutron stars considering the full effect of general relativity. We compare the theoretical values of these quantities with their values inferred from {\\it EXOSAT} data for four low mass X--ray binary sources: XB 1820-30, GX 17+2, GX 9+1 and GX 349+2 and constrain the values of several properties of these sources. According to our calculations, the neutron stars in GX 9+1 and GX 349+2 are rapidly rotating and stiffer equations of state are unfavoured. ", + "introduction": "A low mass X--ray binary (LMXB) is believed to contain either a weakly magnetised neutron star or a black hole as the central accretor. The X--ray emission arises from the innermost region of the accretion disc around the compact star. In the case of a neutron star, there is an additional X--ray component coming from the boundary layer of the star. Mitsuda et al. (1984) showed that the spectrum of a luminous LMXB can be fitted by the sum of a single temperature blackbody spectrum (believed to come from the boundary layer) and a multicolour blackbody spectrum (may be originated from the accretion disc). However these authors used Newtonian models to fit the observed spectra. But near the surface of a neutron star, the accretion flow is expected to be governed by the laws of general relativity due to the presence of strong gravity. Therefore general relativistic models should be used for the purpose of fitting to get the correct best-fit values of the parameters. Besides, the principal motivation behind the study of the spectral and temporal behaviours of neutron star LMXBs is to understand the properties of very high $(\\sim 10^{15}$ g cm$^{-3})$ density matter at the neutron star core (van der Klis 2000). This is a fundamental problem of physics, which can not be addressed by any kind of laboratory experiment. The only way to answer this question is to assume an equation of state (EOS) model for the neutron star core, to calculate the structure parameters of the neutron star and hence to calculate an appropriate spectral model. By fitting such models (for different chosen EOSs) to the observed data, one can hope to constrain the existing EOS models and hence to understand the properties of high density matter. However, general relativistic calculation is essential to calculate the structure parameters of a neutron star and therefore to constrain the EOS models. It is expected that the neutron stars in LMXBs are rapidly rotating due to accretion-induced angular momentum transfer. LMXBs are thought to be the progenitors of milli-second (ms) radio pulsars (Bhattacharya \\& van den Heuvel 1991) like PSR 1937+21 with $P \\sim 1.56$~ms (Backer et al. 1982). The recent discovery of ms $(P \\sim 2.49$~ms) X--ray pulsations in XTE J1808-369 (Wijnands \\& van der Klis 1998) has strengthened this hypothesis. Therefore it is necessary to calculate the structure of a rotating neutron star considering the full effect of general relativity. This was done by Cook, Shapiro \\& Teukolsky (1994) and the same procedure was used by Thampan \\& Datta (1998), to calculate the luminosities of the disc and the boundary layer. The disc temperature profile for a rapidly rotating neutron star was first calculated by Bhattacharyya et al. (2000). These authors also compared their theoretical results with the {\\it EXOSAT} data (analysed by White, Stella \\& Parmar 1988) to constrain different properties of the LMXB source Cygnus X-2. The present work is a continuation of theirs, in which we constrain several properties of four LMXB sources: XB 1820-30, GX 17+2, GX 9+1 and GX 349+2, using the same procedure. These sources were also observed by {\\it EXOSAT} and the data were analysed by White et al. (1988). XB 1820-30 is an atoll source which shows type I X--ray bursts. GX 17+2 and GX 349+2 are Z sources, of which the former shows X--ray bursts. GX 9+1 is an atoll source. As all of them are LMXBs (van Paradijs 1995), the magnetic field of the neutron stars are believed to be decayed to lower values $(\\sim 10^8$ G; see Bhattacharya \\& Datta 1996 and Bhattacharya \\& van den Heuvel 1991). Therefore, we ignore the effect of the magnetic field on the accretion disc structure in our calculations. In section 2, we give the formalism of the work. We present the results and discussion in section 3 and give a summary in section 4. ", + "conclusions": "In this paper, we calculate gravitational mass sequences for different EOS models and constrain several properties of four LMXB sources. For the neutron star in each of the sources, we assume $M = 1.4$~\\msun (i.e., the canonical mass value for neutron stars). We take two values for $\\cos i$ $(i$ is the inclination angle of the source) for each source, namely, $0.2$ and $0.8$. These two widely different values ensure the sufficient generality of our results. For the four sources, the best-fit values (White et al. 1988) of the parameters $T_{\\rm col}^{\\rm max}$, $L_{\\rm D}$ and $L_{\\rm BL}$ are given in Table 1. We take the distance $(D)$ of the source XB 1820-30 as 6.4 kpc (Bloser et al. 2000). We assume $D = 8$~kpc for both GX 17+2 and GX 9+1, as their locations are believed to be near the galactic centre (Deutsch et al. 1999; Hertz et al. 1990) and distance of the galactic centre is $7.9 \\pm 0.3$ kpc, as concluded by McNamara et al. (2000). For GX 349+2, we take $D = 9$~kpc (Deutsch et al. 1999). We display the constrained values with the help of four tables. It is to be noted that here $\\dot M$ is presented in unit of $\\dot M_{\\rm e} = 1.4\\times 10^{17} M/$\\msun~g~s$^{-1}$. The Eddington rate is $\\dot M_{\\rm e}/\\eta$, with $\\eta = E_{\\rm BL}+E_{\\rm D}$. Therefore, as the actual value of $\\eta$ is much lesser than 1.0 (generally not greater than 0.3 and for rapidly rotating neutron star, typically less than 0.2), the value of Eddington accretion rate is much higher than $\\dot M_{\\rm e}$. For all the sources, as the stiffness of the EOS models increases, the absolute values of the allowed spin frequencies $(\\nu_*)$ and rotational frequencies in the ISCO $(\\nu_{\\rm in})$ decreases. This is because, for a stiffer EOS model, neutron star radius is higher and it can support lesser amount of rotation. The energy conversion efficiency is also lesser for a stiffer EOS model (as the neutron star for this case is less compact) and therefore higher accretion rate is needed to generate the observed luminosity (as seen from the tables). In the following, we describe the results for four sources in four subsections and give a general discussion in the last subsection. \\subsection{XB 1820-30} We display the allowed ranges of different parameters for the source XB 1820-30 in Table 2. It is seen that for $\\cos i = 0.2$, the spin frequency $(\\nu_*)$ of the neutron star comes out to be very high. But in the case of $\\cos i = 0.8$, it is not possible to constrain $(\\nu_*)$ for (20\\%,50\\%) uncertainty set (for all EOS models) and for both the uncertainty sets (for EOS model D). The ranges of the colour factor are in general consistent with the results of Shimura \\& Takahara $(f \\sim 1.7 - 2.0)$ and Borozdin et al. (1999) $(f = 2.6)$. However, some discrepancy can be noted with the latter one for softer EOS models \\& $\\cos i = 0.2$. The value of the rotational frequency in the ISCO $(\\nu_{\\rm in})$ comes out to be $\\sim 1$ kHz for all the cases. The values of the stellar equatorial radius are in the range $8 - 21$ km. The peak of the disk effective temperature occurs in the radial range $18-30$ km, and always well-outside (by several kilometers) the neutron star's surface. This shows the validity of eq. (4), as discussed in section 2. The overall range of the accretion rate comes out to be $0.5-31.4 \\dot M_{\\rm e}$. \\subsection{GX 17+2} Table 3 shows the results for the source GX 17+2. Here the ranges of $(\\nu_*)$ are similar to those for XB 1820-30. But for GX 17+2, the value of $i$ is expected to be moderately high (Titarchuk et al. 2001) and therefore $\\cos i$ is not possibly as high as 0.8. It is, therefore, quite likely that the neutron star in this source is rapidly rotating, although no decisive statement can be made. The allowed values for $f$ for GX 17+2 is systematically lower than those for XB 1820-30, and in the case of softer EOS models \\& $\\cos i = 0.2$ they do not tally with the result of Borozdin et al. (1999). The allowed values of $(\\nu_{\\rm in})$ coms out to be $\\sim 1$ kHz, but for EOS model (A) \\& $\\cos i = 0.8$, 2 kHz value is also possible. The ranges of $R$ and $r^{\\rm max}_{\\rm eff}$ are similar to those for XB 1820-30 and the allowed values of the accretion rate are in the range $2.0 - 131.0 \\dot M_{\\rm e}$. \\subsection{GX 9+1} The results for the source GX 9+1 are given in Table 4. For this source, $\\nu_*$-value comes out to be very high for all the EOS models and for both the $\\cos i$-values. Here, the allowed values for $f$ are inconsistent with Borozdin et al. (1999) for softer EOS models \\& $\\cos i = 0.2$. The allowed values of $\\nu_{\\rm in}$ are $\\sim 1$ kHz and the allowed ranges of $R$ and $r^{\\rm max}_{\\rm eff}$ are $10-21$ km and $18-30$ km respectively. The allowed values of accretion rate for this source come out to be in the range $3.8-116.7 \\dot M_{\\rm e}$. \\subsection{GX 349+2} The allowed ranges of different parameters for the source GX 349+2 are given in Table 5. As is the case for GX 9+1, here also the value of $\\nu_*$ comes out to be very high for all the chosen EOS models and $\\cos i$-values. The allowed values for $f$ in general tally with the results of Shimura \\& Takahara (1995), but like other three sources do not match with the result of Borozdin et al. (1999) for softer EOS models \\& $\\cos i = 0.2$. Here $\\nu_{\\rm in}$ comes out to be $\\sim 1$ kHz, and $R$ \\& $r^{\\rm max}_{\\rm eff}$ are in the ranges $10-21$ km and $18-30$ km respectively (like GX 9+1). The accretion rate for this source is in the range $4.5-168.5 \\dot M_{\\rm e}$. \\subsection{General Discussion} Here we have constrained the values of several properties of four LMXB sources. For all of them, the accretion rates come out to be very high (always $\\ge 0.5 ~\\dot M_{\\rm e})$. This is in accord with the fact that these are very luminous sources. The rotation rate of neutron star in each of the sources is very high (close to the mass shed limit) for $\\cos i = 0.2$. This is because, the values of $L_{\\rm BL}/L_{\\rm D}$ are very low for these cases (see Thampan \\& Datta 1998, Bhattacharyya et al. 2000). But, for $\\cos i = 0.8$, rotation rate can not be constrained effectively for the sources XB 1820-30 and GX 17+2. Therefore, for these two sources, no general conclusion (about the values of $\\Omega_{\\rm *})$ can be drawn. However, the allowed ranges (combined for all the cases considered in a table) of $\\Omega_{\\rm *}/\\Omega_{\\rm *,\\mbox{ms}}$ are $0.93-1.00$ and $0.75-1.00$ for the other two sources GX 9+1 and GX 349+2 respectively (here $\\Omega_{\\rm *,\\mbox{ms}}$ is the $\\Omega_{\\rm *}$ at the mass shed limit; see Bhattacharyya et al. 2000 for the mass shed limit values). Therefore the neutron stars in these two sources can be concluded to be rapidly rotating in general. Our calculated allowed ranges for $f$ are in accord with the results obtained by Shimura \\& Takahara (1995). However, if we take the value $f = 2.6$ (obtained by Borozdin et al. 1999), one would require a very stiff EOS model or a mass greater than $M = 1.4$~\\msun for most of the cases with $\\cos i = 0.2$. High frequency quasi--periodic--oscillations (kHz QPO) have been observed for three (XB 1820-30, GX 17+2 and GX 349+2) of the chosen sources. The observed maximum kHz QPO frequencies are 1.100 kHz (XB 1820-30), 1.080 kHz (GX 17+2) and 1.020 kHz (GX 349+2) (van der Klis 2000). Now, as pointed out in Bhattacharyya et al. (2000), the maximum possible frequency (i.e., the shortest time scale) of the system should be given by the rotational frequency in ISCO $(\\nu_{\\rm in}$; col. 5 of the tables). Therefore, the stiffest EOS model D is unfavoured for $\\cos i = 0.2$ for the source XB 1820-30, as the maximum value of $\\nu_{\\rm in}$ $(= 0.941$~kHz, Table 2) is less than the observed maximum kHz QPO frequency. For the same reason, EOS model D is unfavoured for $\\cos i = 0.2$ for the source GX 17+2. It can also be seen from Table 3 that if we use only the narrower limits on the luminosities and colour temperature, EOS model D (for $\\cos i = 0.8)$ and EOS model C (for $\\cos i = 0.2)$ are unfavoured for the same source. Same is true for EOS model D for the source GX 349+2. As we also see from Table 5, EOS model C is unfavoured for $\\cos i = 0.2$ for this source. Therefore, we may conclude that the stiffer EOS models are unfavoured by our results. We have ignored the magnetic fields of the neutron stars in our calculations. Therefore, the necessary condition for the validity of our results is that the Alfv\\'{e}n radius $(r_{\\rm A})$ be less than the radius of the inner edge of the disc. This condition will always be valid if $R > r_{\\rm A}$ holds. Here $r_{\\rm A}$ is given by (Shapiro \\& Teukolsky 1983), \\begin{eqnarray} r_{\\rm A} & = & 2.9 \\times 10^8 {({\\dot {M}\\over \\dot {M}_{\\rm e}})}^{-2/7} \\mu_{30}^{4/7} ({M\\over M_\\odot})^{-3/7} \\end{eqnarray} \\noindent where $M$ is the mass of the neutron star, $\\mu_{30}$ is the magnetic moment in units of $10^{30}$ G cm$^3$ and $r_{\\rm A}$ is in cm. With typical values of the parameters for the chosen sources $(R = 10$~km, $M = 1.4 M_\\odot$ and $\\dot{M} = 10 \\dot {M}_{\\rm e})$, the upper limit of the neutron star surface magnetic field comes out to be about $2 \\times 10^{8}$~G. Therefore, our results are in general valid for the neutron star magnetic field upto of the order of $10^{8}$~G. This is a reasonable value for the magnetic field of neutron stars in LMXBs, as mentioned in section 1. It is also to be noted that our results are valid for a thin blackbody disc. However, as the spectra of the sources were well-fitted by a multicolour blackbody (plus a blackbody, presumably coming from the boundary layer; White et al. 1988), the assumption of thin blackbody disc may be correct." + }, + "0112/astro-ph0112208_arXiv.txt": { + "abstract": "We present ultraviolet images of 27 3CR radio galaxies with redshifts $z<$0.1 that have been imaged with the {\\em Space Telescope Imaging Spectrograph (STIS)} on board the {\\em Hubble Space Telescope (HST)}. The observations employed the NUV-MAMA and broad-band filters with peak sensitivity at 2200\\AA . We find that the UV luminosities show approximately a factor of 10 to 100 higher dispersion than the optical. We compare the UV morphologies with optical $V$- and $R$-band WFPC2 snapshot survey images. We have found dramatic, complex and extended ultraviolet emission from radio galaxies {\\it even at zero redshift}. We find a diverse range of UV morphologies, some completely divergent from their visual morphology, which are reminiscent of the chaotic high-$z$ radio galaxies structures seen in rest-frame UV. The UV morphologies show regions of star formation, jets, and possible scattered AGN continuum. The UV emission is generally not aligned with the radio structure. We also detect the diffuse UV emission of the host galaxy. We propose that these are the same physical phenomena as observed at high redshift, but on a smaller spatial scale. ", + "introduction": "The study of radio galaxies impacts many areas of astrophysics and cosmology. Typically the hosts of powerful radio sources are massive, early-type elliptical galaxies that often lie at the heart of clusters \\citep{mat64,yat89,zir96,hill91}. How these galaxies and clusters are assembled in the early universe from the initial fluctuations is largely unknown. The detailed physics of mergers, competing cooling and heating processes and gravitational aggregation that leads to the massive clusters and giant dominant central galaxies are currently poorly understood. Radio galaxies can be seen at all redshifts, thus providing an important, consistent probe whose systematic characteristics may be traced from the earliest observable times to the present. Not only that, but the fact that they are massive evolving galaxies, often in a most privileged position at the sites of cosmological structure formation, makes them especially important in achieving an accurate picture of the physics and evolution of the universe. While the relationship between AGN evolution and galaxy evolution is unknown, there may be links between radio galaxy activity and star formation \\citep{dey97,bic2000} . The epoch of intense star formation in the universe is similar to the period when the volume density of quasars and radio galaxies was orders of magnitude higher than it is today. Recent X-ray observations show directly that the radio sources themselves interact perhaps in a fundamental way with the hot coronal intra-cluster gas \\citep{smi2001,hard2001} that represents one of the most massive cluster components after the dark matter. Black hole searches have shown that almost every large elliptical galaxy harbors a massive black hole \\citep{mag98,fer2000}, implying all such galaxies go through an active phase. The ubiquity of activity in galaxies thus makes the study of radio galaxies crucial for galaxy evolution and understanding how this phase affects the star formation, gas and dust content, and dynamics of the host galaxy. Investigation of the nuclear environments of radio galaxies is an important area of AGN physics research, with implications for the triggering, fueling and evolution of the active nucleus. At high redshift ($z>$0.6) the UV rest frame continua of radio galaxies are closely aligned with their radio sources \\citep{mcc87,cha87}. The alignment effect implies that the radio source may play a fundamental role in the evolution of the galaxy and clusters of galaxies. The orientations of lower redshift and lower luminosity radio galaxies have been studied in detail with claims of both minor and major axis alignments with respect to the radio lobes \\citep{bau89}. These processes have been well studied at optical wavelengths via the 3CR Imaging Snapshot Surveys conducted with {\\em Hubble Space Telescope} (HST) in Cycles 4-8. Nearly all the extragalactic radio sources in the 3CR catalog \\citep{ben62a,ben62b,spi85} were observed in this series of programs. Impressive results concerning the prevalence of dust disks and optical jets are described in \\cite{mar98,mar99,leh99,mcc97,dek96}. Here we describe an observational program to obtain high spatial resolution HST UV images of nearby powerful 3CR radio galaxies. We present the results of UV imaging of the low redshift ($z<0.1$) subset of the sample. These UV images are a major enhancement to the existing database, and represent the first systematic high spatial resolution survey of radio galaxies in this unexplored spectral window. In particular, using UV imaging we seek to distinguish the relative roles of gas, dust, jets and star formation at the present time. The UV band is well suited to addressing these questions because of the high sensitivity to the youngest, hottest stars, and blue synchrotron jets. Also the UV is optimal for detecting scattered emission because of the blue colour of the illuminating nuclei, and the high efficiency of scattering at shorter wavelengths. A primary motivation for UV imaging of these nearby 3CR sources is to provide a zero redshift comparison sample for the extraordinary rest frame UV morphologies found at high redshift. In any systematic study spanning a significant fraction of the age of the universe it is difficult to disentangle the effects of evolution, intrinsic power and wavelength of the observation. For the first time we can empirically characterize the rest frame UV structure of radio galaxies at zero redshift. The details of the observations and data processing are described in \\S~\\ref{obs_section}. The UV fluxes and extinction-corrected luminosities of all the objects are presented in \\S~\\ref{results_section} along with descriptions of the individual sources and their corresponding optical and radio properties. In \\S~\\ref{analysis_section} the UV and optical luminosities are compared to predictions from star formation and and ionized gas emission models. Discussion is presented in \\S~\\ref{discussion_section} ", + "conclusions": "In this section we summarize the diverse range of UV morphologies and emission structures displayed in Figures~\\ref{3c29} to \\ref{3c465}. Some objects exhibit similar features which allows us to perform a rough categorization into six groups based on the morphology and scale of the UV emission. Figure~\\ref{sum_fig} shows the categorized objects where each row of images comprises a single category. Figure~\\ref{sum_fig} is arranged such that the first four rows display the groups with extended UV emission in order of increasing physical size. \\subsubsection{Dust Disks} The objects which display a dust-disk morphology (3C~40, 3C~270, 3C~296, 3C~449 and 3C~465) define a remarkably uniform category in terms of the scale and physical structure of the of the UV emitting region. These objects are shown in Figure~\\ref{sum_fig} (top row) on the same physical scale (2.0 kpc box). The three most similar of these objects, 3C~40, 3C~270 and 3C~296, show UV emission that unambiguously comes from the smooth underlying elliptical host galaxy, with the dust disk apparent in absorption against the background UV galaxy light. The other two objects in this group, 3C~449 and 3C~465 clearly have similar scale disk structures, and include contribution from the host galaxy starlight, however we note that the detailed UV morphologies of these two dust disks differ from 3C~40, 3C~270, and 3C~296. In the case of 3C~449 the dust disk structure is seen both in emission and absorption. In 3C~465 the disk is not as regular as 3C~40, 3C~270, or 3C~296 and the source contains a compact central core. We also note that 3C~40 is the only object in this group with no compact core seen in the optical image. \\subsubsection{Dust Associated Star Formation} Another striking group of objects are those which exhibit bright extended UV emission over scales of 5 to 20 kpc. In a number of cases this large scale emission is clearly associated with the galaxy dust lane, and in all cases where this is true, the dust lane displays a very unsettled and chaotic morphology (3C~285, 3C~293, 3C~236 and 3C~321 ). In all these cases the UV emission appears to be star formation. This group of objects is displayed with a uniform box size of 20~kpc in the third row of Figure~\\ref{sum_fig}. We note that the UV morphology in these objects is comparable to the rest-frame UV structures seen in optical images of HZRG, and return to this result in the discussion section. \\subsubsection{Complex Extended UV Emission} The objects 3C~405, 3C~305, 3C~192 and 3C~198 also display bright, extended and interesting UV morphology. The emission is typically extended over a 5~kpc scale in these objects. While these objects display quite similar features to the {\\it dust associated star formation} group, we assign them to a separate group because the UV emission is not so clearly identified as star formation, and these objects may represent the presence other UV emission mechanisms. For example, in 3C~405 the bright UV clumps are most likely due to the young star forming regions \\citep{jac98}, but the geometry and presence of strong optical emission lines leads us to expect some contribution from for ionized gas and scattered UV emission. The UV emission in 3C~305 is closely related to the dust lane and may be star formation, however the more diffuse distribution of the UV emission makes this uncertain, and scattered light from the jet may play a role. 3C~192 is smaller, but morphologically similar to 3C~305 and may also be dominated by scattered light. We loosely categorize these objects as {\\it complex extended} UV emission. The UV images of the objects in this category are displayed in the fourth row of Figure~\\ref{sum_fig} with a box size of 10~kpc. \\subsubsection{Compact Core plus Galaxy or Jet Components} A number of objects show compact cores, in combination with a significant contribution from the underlying elliptical galaxy starlight (eg. 3C~310, 3C~35). Some of these objects also display extra UV emission features such as the prominent UV jet seen in 3C~66B, and the possible weak jet in 3C~388. Other objects with compact cores also exhibit absorption of the UV light by a relatively smooth dust lane as in 3C~326, or extra UV emitting sources such as the filament in 3C~317, and the additional UV point source in 3C~338. We group all these objects into a category called {\\it compact core plus galaxy or jet components}. Clearly this is a less homogeneous group of objects, which represent a range of physical mechanisms. The UV images of these sources are displayed in the second row of Figure~\\ref{sum_fig} with a uniform box size of 5.0~kpc. \\subsubsection{Nuclear Dominated} Three of the objects in our sample 3C~227, 3C~382 and 3C~390.3, have very bright UV nuclei at their cores. All of these bright nuclei are all unresolved at the resolution of HST. The UV images of this {\\it nuclear dominated} category are displayed in the fifth row of Figure~\\ref{sum_fig} with a box size of 10~kpc. \\subsubsection{UV Host Galaxy and Weak Detections} The remaining objects consist of the weakly detected UV emission in 3C~353, and the very faint nucleus plus galaxy host emission detected in 3C~29. We display this category in the last row of Figure~\\ref{sum_fig} with a box size of 10~kpc. \\label{discussion_section} As described in the above classifications of our UV images, we find that the processes of star formation, nuclear emission, scattering, and extinction by dust, are the basic components which determine the UV appearance of low redshift radio galaxies. The relationships between these various components provide clues to the physics, merger history and evolution of radio galaxies. These same processes are also the likely ingredients which determine the appearance of HZRG. We defer detailed analysis of the multiwavelength properties and of the individual objects to future papers. Here we briefly discuss some implications of the star forming regions found in the 3C radio galaxies, and make a broad comparison of the low redshift UV morphologies with high redshift radio galaxies. We find a progression of UV morphologies from highly chaotic dusty systems with strong star formation, to smooth regular systems with nuclear dust-disks that are aligned perpendicularly to the radio jet. This range of morphologies appeals to a scenario where the central regions of radio galaxy hosts may evolve from highly chaotic systems, to more stable regular systems. This idea is linked with the merger scenario for the triggering and fueling of the active nucleus. In this evolutionary scenario the categories of objects identified in this paper may roughly represent the various stages of the effect of infalling gas and dust. For example the large scale chaotic dust lanes of the {\\it dust associated star formation} group may represent the earliest stages of a merger event. The radio axes of these objects bear no relation to the dust lanes or central structures in these galaxies, perhaps indicating that these systems have not had enough time to settle into a stable configuration. At the other extreme the {\\it dust disk} objects do not show strong evidence for star formation associated with the dust, and these regular dust disks may represent the final stable configuration. In all the dust disks the radio jet is aligned roughly perpendicularly to the major axis of the disk. One of the difficulties with such a simplified scenario is that the timescales for settling of gas is of order 1~Gy after the merger \\citep{gunn79,chr93}, which is longer than the $\\sim10^7$ year duration of a single epoch of radio activity \\citep{ale87}. If a merger resulted in only a single epoch of radio activity, then there would not be enough time for evolution from a 3C~293 like object into a settled dust disk within the lifetime of the radio jets. The evolutionary scenario may however be credible if the onset of radio activity occured (or re-occurred) at a time after the merger, perhaps dependent on the precise parameters of the encounter. These parameters may include merger impact geometry and gas infall timescale, also the gas and dust content of the merging systems, and the angular momentum loss rate. If the onset of radio activity occurs at random time during a merger we may be able to account for the coexistence of radio jets and the wide range of morphologies as described here. Alternatively FR-I sources may be intriniscally different due to a different low accretion fuelling rate for the AGN as described in \\cite{bau95}. Estimating the ages of the newly triggered star formation may provide a means of dating major merger events in the history of the host galaxy. Similarly, if star formation is related to the nuclear activity of the galaxy, the star formation ages will significantly aid our understanding of these systems as has recently been demonstrated by \\cite{odea01} for the case of 3C~236 (using the UV and optical data described in this paper). They showed that the young star forming regions may be related to the inner 2~kpc scale radio source, and that the star formation may reflect the radio source history. Clearly many of the {\\it dust associated star formation} objects show similarities to 3C~236 and we expect that more accurate dating of the individual star forming knots will shed light on the merger and activity history of these galaxies. A primary motivation for UV imaging of low redshift radio galaxies is to provide a zero redshift comparison sample for the extraordinary rest frame UV morphologies found at high redshift. High resolution HST images of high redshift radio galaxies have been described by \\cite{bes97} ( 28 3CR radio galaxies $0.6 0.6$. ", + "introduction": "\\label{sec:intro} The Wisconsin \\ha\\ Mapper (WHAM) is a dual-etalon, Fabry-Perot spectrometer designed to measure optical emission from diffuse interstellar sources. WHAM is approximately 100 times more sensitive than the previous instrument of its kind, making it possible to study very low levels of emission not accessible in the past. The WHAM \\ha\\ studies presented in this paper were made along three lines of sight in which the interstellar \\ha\\ emission is extremely faint. In particular, we have measured the \\ha\\ emission line profile toward HD~93521 ($\\ell~=~183\\fdg1$, $b~=~+62\\fdg2$), supplementing the work of \\cite{sf93}, who obtained ultraviolet \\emph{HST} absorption line measurements for numerous species toward this star. Secondly, we have measured the \\ha\\ emission toward $\\ell~=~148\\fdg5$, $b~=~+53\\fdg0$, which lies in the region of sky known as the Lockman Window. This region has the lowest \\hi\\ column density of any section of the northern sky, and thorough 21~cm studies of the area have been made by \\cite{jahoda90} and \\cite{lockman86}. The \\ha\\ observations provide information about the amount of H~II in this unusual direction. Lastly, we have observed the line of sight toward $\\ell~=~163\\fdg5$, $b~=~+53\\fdg5$ (referred to below as ``Off A''). This direction was selected for investigation based on WHAM Northern Sky Survey data, which indicated that only minimal levels of \\ha\\ emission could be present. Results for these three directions are discussed in \\S\\ref{sec:ha_93521}, \\S\\ref{sec:lm}, and \\S\\ref{sec:offA}, respectively. Accompanying WHAM's ability to detect faint \\emph{Galactic} emission is its ability to detect faint \\emph{terrestrial} emission lines as well. In fact, WHAM \\ha\\ spectra centered near the local standard of rest (LSR) show a total of twelve weak ($\\sim$0.03 -- 0.15~R, 1~R~=~10$^{6}$/4$\\pi$ photons cm$^{-2}$ s$^{-1}$ sr$^{-1}$) sky lines of unknown origin in addition to the well known and much brighter geocoronal \\ha\\ line. These weak lines are present at fixed positions with respect to the geocoronal line, and their intensities vary slightly during the night. Of these twelve lines, four on the blue side of the geocoronal line were first noticed in WHAM spectra used to study high-velocity clouds and are discussed by \\cite{tufte98} and by \\cite{tufte97}. One of the lines to the red of the geocoronal line is discussed by \\cite{haffner98}. Since the atmospheric lines can significantly affect measurements of faint Galactic \\ha\\ emission, it is important to understand and remove their contributions to the data. In \\S\\ref{sec:atmlines}, we give a characterization of the twelve faint atmospheric lines based on a study of a large sample of WHAM \\ha\\ spectra. In addition to \\ha\\, other, fainter emission lines can provide information about the physical state of the emitting gas. For example, measurements of Galactic \\oiw\\ emission along a given line of sight, when combined with the \\ha\\ intensity for that direction, yield information regarding the ionization state of the emitting gas, a subject of interest and some controversy for the HD 93521 sightline (e.g., \\citealt{sf93, smh00}). In particular, the hydrogen ionization ratio $n(\\mathrm{H^+})$/$n(\\mathrm{H^\\circ})$ in the gas can be obtained via the observed \\oi/\\ha\\ line intensity ratio. The precise relation, as given by \\cite{reynolds98} and references therein, is \\begin{equation} \\frac{I_\\mathrm{[O~I]}}{I_{\\mathrm{H}\\alpha}} = 2.74 \\times 10^4 \\left(\\frac{T_4^{1.854}}{1 + 0.605T_4^{1.105}}\\right) \\mathrm{exp}\\left( -\\frac{2.284}{T_4} \\right) \\frac{n(\\mathrm{O})}{n(\\mathrm{H})} \\left[ \\frac{1 + n(\\mathrm{H^{\\circ}})/n(\\mathrm{H^+})} {1 + (8/9)n(\\mathrm{H^{+}})/n(\\mathrm{H^\\circ})} \\right], \\end{equation} where $I_\\mathrm{[O~I]}/I_{\\mathrm{H}\\alpha}$ is the intensity ratio measured in energy units and $T_4$ is the electron temperature in units of 10$^4$~K. The term \\emph{n}(O)/\\emph{n}(H) is the gas-phase abundance of oxygen, which we take to be $3.2~\\times~10^{-4}$ \\citep{mjc98}. The [O~I] observations are discussed in \\S\\ref{sec:oi_93521} and \\S\\ref{sec:disc}. ", + "conclusions": "\\label{sec:disc} Interstellar \\ha\\ emission has been detected and characterized toward two high Galactic latitude, low H~I column density sightlines, the Lockman Window ($\\ell~=~148\\fdg5$, $b~=~+53\\fdg0$) and HD 93521 ($\\ell~=~183\\fdg1$, $b~=~+62\\fdg2$). The relatively bright geocoronal line and the presence of many lower intensity, unidentified atmospheric emission lines make it difficult to measure accurately the interstellar \\ha\\ emission in these faintest parts of the sky. However, from careful, long integration WHAM spectra spread over many months, the interstellar and terrestrial emissions were successfully separated. The results provide new constraints on the nature of the interstellar clouds along these two sightlines and their environment. We have found that the Lockman Window, in addition to having the lowest H~I column density in the sky ($5 \\times 10^{19}$ cm$^{-2}$), also has unusually weak interstellar \\ha\\ emission (0.20~R), well below the 0.8~R average for this Galactic latitude (from the WHAM sky survey; \\citealt{survey}). This region of the sky thus appears to be a true low column density window through the Galactic disk, depleted both in H~I and in H~II. Only one \\ha\\ velocity component is detected toward the Lockman Window, the component near the LSR, even though in the H~I spectrum there is also a prominent emission component near --50~\\kms\\ (see Fig.~\\ref{fig:main}). An \\ha\\ intensity of 0.20~R for the Lockman Window implies an emission measure of 0.55 cm$^{-6}$ pc (at 10$^4$ K), which corresponds to a column density N$_{\\mathrm{H~II}} \\approx 2 \\times 10^{19}$ cm$^{-2}$, if the mean density within the ionized regions is about 0.08 cm$^{-3}$ (see \\citealt{reynolds91}). This is approximately 40\\% the H~I column density and implies that the total hydrogen column density in this direction is about $7 \\times 10^{19}$ cm$^{-2}$. Toward HD 93521, $21\\arcdeg$ away, the slow and intermediate velocity components detected in \\ha\\ clearly correspond to the two principal emission components present in the 21 cm spectrum. The intensity of the intermediate velocity (--50~\\kms) component relative to the slow (--10~\\kms) component is significantly weaker in the \\ha\\ spectrum than in the 21~cm spectrum, suggesting that perhaps an intermediate velocity \\ha\\ component toward the Lockman Window is just below the detection limit of the observation. For HD 93521 there is also strong evidence for a high velocity \\ha\\ emission component at --90~\\kms, but no corresponding feature in the H~I spectrum. This sightline is very near the High Velocity H~I Cloud Complex M, parts of which have been detected in \\ha\\ at radial velocities ranging from --109~\\kms\\ to --61~\\kms\\ \\citep{tufte98}. Therefore, the high velocity \\ha\\ component may be associated with a fully ionized portion of this high velocity cloud complex. The total \\ha\\ intensity (0.37~R) for the HD 93521 sightline corresponds to an emission measure of 1.0 cm$^{-6}$ pc, and thus N$_{\\mathrm{H~II}} \\approx 4 \\times 10^{19}$ cm$^{-2}$ (about 30\\% of the H~I column density of $1.25 \\times 10^{20}$ cm$^{-2}$). Although there is good correspondence between the H~II and H~I for the low velocity and (for HD 93521) the intermediate velocity components, the difference in the general appearance of the \\ha\\ and 21 cm line profiles strongly suggests that the H~II and the H~I are not mixed together in the form of partially ionized clouds. In particular, the \\ha\\ components appear systematically broader. For example, the width (FWHM) of the intermediate velocity H~I component toward HD 93521 is 20 \\kms, only about half that of the corresponding \\ha\\ emission. As a result, there is significant blending between the slow and intermediate velocity components in the \\ha\\ spectrum, whereas these components are well resolved in the H~I spectrum. This is not a result of the different spectral resolutions of the \\ha\\ and 21~cm observations. Therefore, rather than being primarily neutral clouds that are 30\\%--40\\% ionized, as proposed by \\cite{sf93} and \\cite{sciama98}, each component seems to consist of separate regions of H~I and H~II at nearly the same radial velocity. This could be the result of either a close physical association between the H~I and H~II \\citep{mo77} or large scale organized motions of the interstellar medium along this sightline \\citep{wc01}. Because the \\oiw/\\ha\\ intensity ratio is a probe of the hydrogen ionization fraction within the emitting gas (see \\S\\ref{sec:intro}), observations of \\oi\\ provide a potentially rigorous test of these ``mixed'' versus ``separated'' models for the H~I and H~II in these clouds. Unfortunately, the \\oi\\ spectrum, which only allows an upper limit on the intermediate velocity (--50~\\kms) component (Table~\\ref{table:oi_results}; see also \\S\\ref{sec:oi_93521}), does not provide an unambiguous result. If the temperature is as low as 6000 K, as adopted by \\cite{sf93} based on the widths of lines generally associated with H~I clouds, then the \\oi\\ observations are consistent with their model of primarily neutral clouds containing a 30\\% mixture of H$^+$. On the other hand, if the temperature is closer to 10,000 K, as recent observations of [S~II] $\\lambda$6717 and [N~II] $\\lambda$6584 suggest (Pifer et al, in preparation), then these \\oi\\ results imply that the emitting regions are predominantly ionized (i.e., $n(\\mathrm{H}^+)/n(\\mathrm{H}_{\\mathrm{total}}) > 0.6$). The \\oi\\ limit is also just consistent with recent predictions of a supernova remnant ionization model for the HD 93521 sightline by \\cite{smh00}. Given the severe atmospheric line contamination of this region (worse than at \\ha ), it will be extremely difficult with ground-based observations to push the sensitivity limit of the \\oi\\ observations down to the level of 0.03 R or lower needed to test these models more definitively. If the clouds in these two sightlines are in fact ionized by an external flux of ionizing radiation, then the \\ha\\ intensity of each H~I cloud is a measure of the incident ionizing flux. Specifically, the one-sided incident flux is $2.2 \\times 10^6$ $I_{\\mathrm{H}\\alpha}$ photons cm$^{-2}$ s$^{-1}$, for $I_{\\mathrm{H}\\alpha}$ in Rayleighs and a temperature of 10$^4$ K (e.g., \\citealt{reynolds95, tufte98}). These observations therefore imply a variation in the intensity of the ambient ionizing flux, from less than $1.3 \\times 10^5$ photons cm$^{-2}$ s$^{-1}$ at the intermediate velocity (--50~\\kms) H~I cloud in the Lockman Window to about $4.4 \\times 10^5$ photons cm$^{-2}$ s$^{-1}$ for the low velocity clouds. The \\ha\\ intensity of the high velocity (--90~\\kms) component toward HD 93521 gives a flux of $5 \\times 10^4$ photons cm$^{-2}$ s$^{-1}$; however, because there is no associated H~I detected at this velocity, the H~II is probably density bounded, and thus the derived value for the ionizing flux must be considered as a lower limit. If the higher velocity clouds are farther from the Galactic midplane than the lower velocity clouds \\citep{kuntz96}, these results indicate a decrease in the ionizing flux with distance above the plane." + }, + "0112/astro-ph0112446_arXiv.txt": { + "abstract": "We review the present status of the Baikal Neutrino Project and present the results obtained with the deep underwater neutrino telescope {\\it NT-200} \\vspace{1pc} ", + "introduction": "The Baikal Neutrino Telescope is deployed in Lake Baikal, Siberia, \\mbox{3.6 km} from shore at a depth of \\mbox{1.1 km}. The present stage of the telescope, {\\it NT-200} \\cite{APP}, was put into operation at April, 1998. Results of searches for atmospheric neutrinos, WIMPs and magnetic monopoles obtained with {\\it NT-200} have been presented elsewhere \\cite{B2001}. During three winter seasons, starting with 1998, a Cherenkov EAS array, consisting of four \\mbox{\\textit{QUASAR--370}} phototubes was deployed on the ice, just above the underwater telescope, with the aim to study the angular resolution of the latter. Analysis of data show that the angular resolution of underwater telescope for vertical muons after modest cuts is about 4$^0$. In the last winter expedition we continued to study the feasibility of acoustic detection of EAS cores in water with an EAS array and four hydrophones. During the EAS array life time of 154 hours, almost 2400 showers with energies above 5 PeV have been recorded. Coincidence data of the EAS array and hydrophones are presently analyzed. Also investigations of water parameters have been continued. Independent measurements of light absorption and scattering have been performed by BAIKAL and NEMO (A.Capone et al.) groups. Preliminary results indicate that the two independent sets of optical data are in a good agreement. In the course of the last expedition we also lowered a special string with instruments for diverse goals, in particular to measure the group velocity of light in water at two different wavelengths and to test a two-channel optical module and a calibration light beacon. Below, we present new results of a search for a diffuse high energy neutrino flux with the neutrino telescope {\\it NT-200}. ", + "conclusions": "The neutrino telescope {\\it NT-200} is taking data since April 1998. It performs investigations of atmospheric muons and neutrinos, and searches for WIMPs, magnetic monopoles and extraterrestrial high energy neutrinos. \\begin{figure}[htb] \\includegraphics[width=5.0cm,height=6.3cm]{enh_nt_col.eps} \\vspace{-0.7cm} \\caption{Sketch of {\\it NT-200}+. } \\end{figure} In the next 2 years we plan to increase the sensitivity to diffuse fluxes by a factor of four. With a moderate upgrade of only 22 optical modules at three additionals strings we would reach a sensitivity of $\\Phi_{\\nu}E^2 \\leq 3\\cdot10^{-7}$cm$^{-2}$s$^{-1}$sr$^{-1}$GeV. This upgrade towards a 10\\,Mton detector {\\it NT-200}+ is sketched in fig.3. This work was supported by the Russian Ministry of Research (contract \\mbox{\\sf 102-11(00)-p}), the German Ministry of Education and Research and the Russian Fund of Basic Research (grants \\mbox{\\sf 99-02-18373a}, \\mbox{\\sf 01-02-31013},\\mbox{\\sf 00-15-96794} and \\mbox{\\sf 01-02-17227}), and by the Russian Federal Program ``Integration'' (project no. 346)." + }, + "0112/astro-ph0112393_arXiv.txt": { + "abstract": "Cosmic shear probes the distribution of dark matter via gravitational lensing of distant, background galaxies. We describe our cosmic shear survey consisting of deep blank fields observed with the Keck II telescope. We have found biases in the standard weak lensing analysis, which are enhanced by the elongated geometry of the Keck fields. We show how these biases can be diagnosed and corrected by masking edges and chip defects. ", + "introduction": " ", + "conclusions": "" + }, + "0112/astro-ph0112116_arXiv.txt": { + "abstract": "We have recently developed the capability to make solar vector (Stokes IQUV) magnetograms using the infrared line of MgI at 12.32\\ \\micron. On 24 April 2001, we obtained a vector magnetic map of solar active region NOAA 9433, fortuitously just prior to the occurrence of an M2 flare. Examination of a sequence of SOHO/MDI magnetograms, and comparison with ground-based H-alpha images, shows that the flare was produced by the cancellation of newly emergent magnetic flux outside of the main sunspot. The very high Zeeman-sensitivity of the 12\\ \\micron\\ data allowed us to measure field strengths on a spatial scale which was not directly resolvable. At the flare trigger site, opposite polarity fields of $2700$ and $1000$ Gauss occurred within a single $2$ arc-sec resolution element, as revealed by two resolved Zeeman splittings in a single spectrum. Our results imply an extremely high horizontal field strength gradient ($5$ G/km) prior to the flare, significantly greater than seen in previous studies. We also find that the magnetic energy of the cancelling fields was more than sufficient to account for the flare's X-ray luminosity. ", + "introduction": "Solar flares are believed to be powered by the release of magnetic energy, but the details of the process are not understood. Flare-related changes in magnetic configurations of active regions have long been sought, without clear success. However, recent observations using ground-based \\citep{cam99,har01} and space-borne magnetographs \\citep{kos99, kos01} are beginning to show clear signatures of magnetic energy release, temporally coincident with large flares. It has long been known that small to moderate flares can occur where newly-emergent flux becomes abutted against its opposite polarity \\citep{mar84} and is cancelled, presumably by reconnection. Flux-cancellation represents one obvious type of flare-related magnetic change, but it is also believed that flux cancellation can trigger the release of magnetic energy stored over larger volumes \\citep{pre84}. The derivation of the magnetic energy from observations requires accurate measurement of the magnetic field strength, and conventional magnetographs can give large errors in field strength. Because Zeeman splitting for visible-region lines is less than the intrinsic line width, it can be difficult to discriminate between a strong field having a small filling factor, and a much weaker field with unit filling factor, and these two cases can represent vastly different magnetic energies. In recent years, infrared (IR) lines have been increasingly used \\citep{meu98}, because they can exhibit resolved Zeeman splitting, giving the field strength independent of filling factor for fields stronger than some limit (typically $\\sim 400$ Gauss at 12\\ \\micron). We have been developing instrumentation to make vector field maps (Stokes $IQUV$) of solar active regions using the far-IR line of MgI at 12.32\\ \\micron. One of our vector field maps was fortuitously made just prior to the occurrence of an M2 flare initiated by flux cancellation. Our data have insufficient temporal resolution to demonstrate a decrease in magnetic energy at the precise time of the flare. Nevertheless, we exploit the great Zeeman-sensitivity of these data to demonstrate that the flare began where the horizontal gradient in field strength was exceptionally large ($\\gtrsim 5$ G/km), and that the magnetic energy in the immediate vicinity of the flux-cancellation site was more than ample compared to the flare's X-ray luminosity. ", + "conclusions": "The 350 Gauss panel of Figure 4 shows a noticeably more `diffuse' appearance than the panels at higher field strengths. This shows that weak fields are present in active regions. They are broadly distributed, in contrast to strong fields - which concentrate at specific locations. A large fraction of the weak field component is due to the sunspot `canopy', which can be seen as the extended region of positive polarity most obvious in Stokes-V on the side of the sunspot toward disk center (compare Figure 2). We have previously described a prominent sunspot canopy observed in the 12.32\\ \\micron\\ line \\citep{jen01}. The sunspot penumbral field shows a ring-like morphology, which of course reflects the fact that the single field strength in each panel is found at only a particular radius from the spot center. With increasing field strength, the width and radius of the ring decrease. At the largest field strength (3900 Gauss), the polarity of the sunspot field seems to reverse. We have observed this effect in several sunspots, and it is certainly not an observational artifact. However, it does not represent an actual polarity reversal. Recall that this emission line is flanked by a broad, shallow, absorption trough \\citep{chasch91}. At large distances from line center, the sunspot Stokes-V profile is dominated by the relatively low-amplitude V signal from the absorption trough, which accounts for the apparent polarity reversal. In principle, this effect conveys information on the height gradient of field strength, since the absorption is formed deeper than the emission. However, a quantitative interpretation is beyond the scope of this paper. The MDI whole-disk magnetogram at 17:36 UT registers peak field strengths of $\\sim 250$ and $\\sim 500$ Gauss for the positive polarity fields in delta regions 1 and 2 respectively. In contrast, the 12\\ \\micron\\ data show that these regions are not even prominent in positive polarity until the field strength exceeds $700$ Gauss, with significant intensity being recorded even at $2800$ Gauss. Evidently the fields are not resolved at the MDI whole disk resolution of $2$ arc-sec. Although our spatial resolution was also $2$ arc-sec, the very large Zeeman splitting in the 12\\ \\micron\\ line permits us to measure true field strengths without spatially resolving the fields. In both delta regions 1 and 2 the average strength of the positive polarity fields is $\\sim 2000$ Gauss, but in delta region 2 (where the flare was triggered), the distribution of field strengths extends to greater values than in region 1. This is noticeable in the 3900 gauss panel, where a trace of delta region 2 can still be discerned. Moreover, these fields are in quite close proximity to the small negative polarity pores, producing exceptionally large field strength gradients. Figure 5 shows the Stokes-V profile at the point in delta region 2 where opposite polarities are in contact (marked by the circle on the 2100 Gauss panel). Within the $2$ arc-sec spatial resolution element, two distinct Zeeman splittings are seen: a 2700 Gauss splitting at positive polarity, and 1000 Gauss at negative polarity. These splittings are marked on Figure 5. We contemplated alternative interpretations of this complex Stokes-V profile, which might lessen the necessity for strong opposite-polarity fields in close spatial coincidence. However, alternate interpretations are difficult to reconcile with the Stokes-V profiles in immediately adjacent regions, which support the picture of strong fields in both polarities. This large field strength difference in such close proximity requires a horizontal field gradient of $\\sim 5$ G/km. To our knowledge, this is the largest horizontal gradient ever observed on the Sun, and suggests reconnection at this site, at or near the temperature-minimum height. In their study of upper-photospheric reconnection, \\citet{litmar99} remark that `the local field cannot be...measured directly.' The high magnetic resolution afforded by the 12\\ \\micron\\ line indeed makes such a measurement possible, and the height of line formation is well matched to the temperature-minimum region where the electrical resistivity is maximum. We therefore point out the significant diagnostic potential of the 12\\ \\micron \\ line in studies of photospheric reconnection. Other investigators find high field strength gradients at flaring sites, but less than we infer here. \\citet{kos01} found a gradient of $1.3$ G/km for the X5.7 `Bastille Day' flare; our results for a significantly weaker flare suggest that conventional magnetographs (even space-borne ones such as SOHO/MDI) underestimate the field strength gradients. If so, they probably also underestimate the magnitude of magnetic energy release inferred from observed changes in the fields. Our results have significant implications for the magnetic energy available locally (i.e., in the flux-cancellation region of the mid- to upper photosphere) to power this flare. The total magnetic energy is proportional to the volume integral of $\\frac{B^{2}}{8\\pi}$ \\citep{chan}. We estimate this quantity directly from the Stokes-V profiles: \\begin{equation} E_{tot} = \\frac{\\delta H}{8\\pi} \\epsilon \\sum B^{2} f_{B} dA \\end{equation} where the volume integral is approximated as a discrete sum over field strengths within a spatial resolution area $dA$, times a layer thickness $\\delta H$. $\\epsilon$ is the fraction of the atmosphere which is magnetic, i.e. a filling factor. We constrain $\\epsilon > 0.15$ by comparing the MDI and 12\\ \\micron\\ field strengths. The factor $f_{B}$ expresses the fraction of the magnetic atmosphere which exhibits field strength $B$. The essential advantage of a strongly-split infrared line is that we can infer $f_{B}$ as proportional to the amplitude of the Stokes-V profile at each splitting (i.e., field strength) value. (In so doing we assume that the line formation mechanism for the 12.32\\ \\micron\\ transition is independent of field strength.) Moreover, because the 12\\ \\micron\\ line is formed in the temperature minimum region, $\\sim400$ km above the photosphere \\citep{cha91, car92}, we know that $\\delta H > 400$ km. On this basis we find that the total `local' magnetic energy of region 1 was $> 3 \\times 10^{22}$ Joules, and of region 2, $> 4.5 \\times 10^{22}$ Joules. Using the single-value MDI field strengths with unit filling factor gives $3.5 \\times 10^{20}$ and $6 \\times 10^{21}$ Joules for regions 1 and 2 respectively. By comparison, the X-ray luminosity of the flare was $3 \\times 10^{21}$ Joules. As is well known, only a fraction of the total magnetic energy is available to a flare \\citep{low90}, and the flare's total luminosity will of course be greater than the energy measured in the GOES X-ray bands (Table 1). Hence the MDI energies are uncomfortably small, but the 12\\ \\micron\\ field strengths indicate an ample local resevoir of magnetic energy. A conventional view is that cancelling magnetic flux simply triggers the release of magnetic energy stored in much larger volumes. While this may actually happen for many of the largest flares, our results imply that sufficient energy was available to power this M2 flare `locally', without invoking the release of magnetic energy from much larger volumes." + }, + "0112/astro-ph0112320_arXiv.txt": { + "abstract": "Recent experimental results from supernovae Ia observations have been interpreted to show that the rate of expansion of the universe is increasing. Other recent experimental results find strong indications that the universe is ``flat.'' In this paper, I investigate some solutions of Einstein's field equations which go smoothly between Schwarzschild's relativistic gravitational solution near a mass concentration to the Friedmann-Lema\\^\\i tre expanding universe solution. In particular, the static, {\\it curved-space extension} of the Lema\\^\\i tre-Schwarzschild solution in vacuum is given. Uniqueness conditions are discussed. One of these metrics preserves the ``cosmological equation.'' We find that when the rate of expansion of the universe is increasing, space is broken up into domains of attraction. Outside a domain of attraction, the expansion of the universe is strong enough to accelerate a test particle away from the domain boundary. I give a {\\it domain-size--mass relationship}. This relationship may very well be important to our understanding of the large scale structure of the universe. ", + "introduction": "Recently de Bernardis {\\it et al.}\\cite{deB} reported that a very careful examination of the cosmic microwave background provides, as further explained by Hu\\cite{Hu}, very strong evidence that the universe is flat! In addition, data has recently been reported by Riess {\\it et al.}\\cite{Riess}, and Perlmutter {\\it et al.}\\cite{Perl} on measurements of the luminosity and the redshift of a number of high-redshift supernovae of type Ia. The authors' best fit to their data involves a strongly curved space-time. The curvature constant $\\Omega _k$ (defined below) that they find is about -1.05, which corresponds to a strongly curved, open universe. However, their error bars are sufficiently large so that a flat universe is not inconsistent with their data. These authors have analyzed their data on the basis of the ``cosmological equation'' between the rate of expansion, the mean energy density, the radius of curvature of space, and the cosmological constant. This relation has been derived from Einstein's field equations\\cite{Peb} under the assumption of the validity of the Friedmann-Lema\\^\\i tre line element at large distances. The authors have concluded, in terms of this model of the universe, that instead of the rate of expansion decreasing, as many workers had thought, their data is best fit by a model in which the rate of expansion is increasing. Riess {\\it et al.}\\cite{Riess} find that the deceleration/acceleration parameter (as defined in section VI) $q_0< 0$ (acceleration) with better than 90\\% confidence, and Perlmutter {\\it et al.}\\cite{Perl} find the same at the 2.86 standard deviation level of confidence. In this paper, I investigate some different metrics which may facilitate the investigation of some of the consequences of the reported acceleration of the expansion of the universe. I will focus on spherically symmetric models with a central mass concentration. The fact that the rate of expansion of the universe is increasing leads to the conclusion that the universe is broken up into domains of attraction. Briefly, the underlying physics of this feature may be seen in the following rough analysis. The equation of purely radial motion of a test particle at rest with respect to the Friedmann-Lema\\^\\i tre coordinate system is just \\begin{equation} \\ddot r =\\frac{\\ddot a}ar, \\label{1.1} \\end{equation} where $r$ is the proper distance, and $a$ is the universal expansion factor which is a function of time alone. The most simple way to approximate the effects of a gravitating mass concentration is just to add Newton's force term so that \\begin{equation} \\ddot r =\\frac{\\ddot a}ar -\\frac {GM}{r^2}. \\label {1.2} \\end{equation} If the expansion of the universe is decelerating, then $\\ddot a < 0$, so $\\ddot r <0$ always. Near a gravitating mass the expansion is clearly unimportant and Newton's laws hold with only minimal corrections. On the other hand, if the expansion is accelerating, then $\\ddot a >0$. Thus there will be a distance such that the expansion of the universe exactly balances the gravitational attraction. Test particles at smaller distances will be accelerated inwards. I call this region a domain of attraction. Test particles outside will be accelerated outwards. Hence in this very simple case there are two attractors, one is the mass concentration and the other is the point at infinity. We give a {\\it domain-size--mass relation} in equations (\\ref{8.3}) and (\\ref{8.6}). In the decelerating case, there is just one attractor and its domain of attraction is the entire space. In the second section we remind the reader of the Lema\\^\\i tre-Tolman formalism. This formalism has the feature that the proper time appears explicitly. Thus space-time is divided into a set of three-dimensional, space-like manifolds which are indexed by the proper time. Since it is the proper time which appears in the universal expansion function in the Friedmann-Lema\\^\\i tre line element and it governs the universe as a whole, it seems important to preserve the property that the universal expansion factor is a function of the proper time alone. In the third section we review three different ways to go from the Schwarzschild metric of a planetary system to expanding space at large distances. The McVittie solution has a problem at the Schwarzschild radius using normal expansion factors. The Einstein-Straus ``Swiss Cheese Model'' is unstable to perturbations and there are orbits which are discontinuous functions of their initial conditions. The Bona-Stela model using Liebovitz metric insertions in a Friedmann-Lema\\^\\i tre background predicts an increase in the length of the earth's year which is in strong disagreement with observations. In the fourth section we consider the idea that the vacuum is something more than just empty space. The self-energy of the vacuum may correspond to a mass-energy density of empty space. We know on the laboratory scale from, for example, the Casimir effect, and on the microscopic scale from particle physics that the vacuum does have measurable effects. So far as I know, particle theorists have not yet been able to compute a quantitative result for the energy density. However, they do feel that such an idea is extremely plausible. We consider in this section a solution to Einstein's field equations which has a homogeneous vacuum, mass-energy density and so preserves the so called ``cosmological equation.'' This solution reduces precisely to the exterior Schwarzschild solution near the central mass condensation. In addition, it reduces to the Friedmann-Lema\\^\\i tre solution far from the mass concentration. This solution is continuously, infinitely differentiable everywhere, except at the mass concentration. In the fifth section, I explore some of the properties of the solution obtained in the fourth section. Both ``flat'' and curved space are considered. The limiting results are as expected and the transition between the two aforementioned limits are illustrated. In particular the static, {\\it curved-space extension} of the Lema\\^\\i tre-Schwarzschild solution is given for a mass concentration in a vacuum. In the sixth section, I consider the implications of the increasing rate of expansion of the universe on some of the large scale structures found in the universe. It is found that this feature creates domains of attraction. A {\\it domain-size--mass relation} is derived. Outside these domains, the increasing rate of the expansion of the universe would, in time, be expected to tear the structures apart. It is suggested, from the correspondence between the predicted size of these domains of attraction, the size of the Local Group, and the size of the Virgo Cluster, that this effect may well be important in any study of the large scale structure of the universe. ", + "conclusions": "" + }, + "0112/gr-qc0112017_arXiv.txt": { + "abstract": "We have computed the eigenfrequencies of $f$ modes for a constant-rest-mass sequences of rapidly rotating relativistic inviscid stars in differential rotation. The frequencies have been calculated neglecting the metric perturbations (the relativistic Cowling approximation) and expressed as a function of the ratio between the rotational kinetic energy and the absolute value of the gravitational energy of the stellar model $\\beta \\equiv T/|W|$. The zeros and the end-points of these sequences mark respectively the onset of the secular instability driven by gravitational radiation-reaction and the maximum value of $\\beta$ at which an equilibrium model exists. In differentially rotating stars the secular stability limits appear at a $\\beta$ larger than those found for uniformly rotating stars. Differential rotation, on the other hand, also allows for the existence of equilibrium models at values of $\\beta$ larger than those for uniformly rotating stars, moving the end-point of the sequences to larger $\\beta$. As a result, for some degrees of differential rotation, the onset of the secular instability for $f$ modes is generally favoured by the presence of differential rotation. ", + "introduction": "Instabilities of rotating relativistic stars have been studied for more than three decades in relation to the emission of gravitational waves driven by radiation-reaction. Since the discovery of the so called Chandrasekhar-Friedman-Schutz (CFS) instability (Chandrasekhar 1970 ; Friedman and Schutz 1978), in fact, the stability properties of a number of non-axisymmetric stellar oscillations have been investigated using different techniques (see Friedman 1998 ; Andersson and Kokkotas 2001 for recent reviews). The $f$ (fundamental) modes were among the first modes of oscillations ever to be studied in rotating relativistic stars. These are spheroidal modes with harmonic indices $l=m$ and represent the generalization of the Kelvin modes of Newtonian Maclaurin spheroids to compressible fluid stars. Over the years, the literature on the subject has been continuously updated and the limits for the onset of the instability have been improved through different approaches. This has been done by several groups focussing on Newtonian stellar models (Bardeen et al. 1977 ; Comins 1979a, 1979b ; Clement 1979 ; Durisen and Imamura 1981 ; Imamura et al. 1985 Managan 1985 ; Ipser and Lindblom 1990) on post-Newtonian stellar models (Cutler and Lindblom 1992) and on rapidly rotating relativistic stars (Yoshida and Eriguchi 1997 ; Stergioulas and Friedman 1998 ; Morsink et al. 1999 ; Yoshida and Eriguchi 1999). Part of the interest in these modes comes from the fact that they have long been regarded as the modes of oscillation most susceptible to the CFS instability. In particular, for stars rotating at a rate at which mass starts being shed at the equator (the mass-shedding limit), general relativistic calculations on the stability limit indicate that the $l=m=2$, $f$-mode (also referred to as the ``bar-mode'') could have the shortest growth timescale. Because of this, and because of a possible weakening of the bulk viscosity at very high temperatures (Lai 2001), the bar-mode may represent the most important non-axisymmetric instability in very hot and rapidly rotating newly born neutron stars. The occurrence of an $f$-mode instability in a newly born uniformly rotating neutron star is in general prevented by the large bulk and shear viscosities of nuclear matter in those conditions and detailed calculations have shown that the instability is suppressed except for very large rotation rates, close to the mass-shedding limit (Ipser and Lindblom 1991); for superfluid stars an even stronger damping was calculated (Lindblom and Mendel 1995). More recently, however, a number of new elements have improved our understanding of the instability and have again increased the expectations that the $f$ mode instability might characterize the earliest life stages of a newly formed neutron star. The first of these new elements was provided by Cutler and Lindblom (1992) who have shown that, within a post-Newtonian approximation, general relativistic effects tend to further destabilize the $f$ mode, lowering the critical value of the ratio between the stellar rotational kinetic energy and the absolute value of the gravitational energy, $\\beta_c \\equiv (T/|W|)_c$ at which the secular $f$-mode instability is triggered. Analogous results have also been found by Yoshida and Eriguchi (1997) within the relativistic Cowling approximation and by Stergioulas and Friedman (1998) in fully general relativistic calculations. The second new element was provided by Shibata and Ury\\={u} (2000) whose fully general relativistic hydrodynamical simulations have shown that the remnants of binary neutron star mergers could be, at least for polytropic equations of state, rapidly and differentially rotating stars. In addition to this, Liu and Lindblom (2001) have recently computed the structure of objects formed in accretion-induced collapse of rotating white dwarfs and found that these objects can rotate extremely rapidly and differentially. Furthermore, a differentially rotating object could be produced also during the process of core-collapse in supernova leading to the formation of a neutron star. Within this framework, differential rotation has two major consequences. Firstly, it allows for the existence of an equilibrium model at values of $\\beta$ which are considerably larger than the ones supported by the counterparts with the same rest-mass but uniform rotation (Baumgarte et al. 2000). Secondly, as shown by a number of authors for Newtonian stellar models (Managan 1985, Imamura et al. 1985, Imamura and Durisen 2001), it increases the critical value $\\beta_c$ for the onset of the secular instability. In this revised picture, we have computed the eigenfrequencies of $f$ modes and determined the secular stability limits for rapidly rotating relativistic stars with differential rotation. ", + "conclusions": "" + }, + "0112/astro-ph0112099_arXiv.txt": { + "abstract": "We apply the method of principal component analysis to a sample of simple stellar population to select some age sensitive spectral indices. Besides the well-known age sensitive index, H$\\beta$, we find some other spectral indices have great potential to determine the age of stellar population, such as G4300, Fe4383, C$_2$4668, and Mg$b$. In addition, we find these spectral indices sensitivity to age depends on the metallicity of SSP, H$\\beta$ and G4300 are more suited to determine the age of low metallicity stellar population, C$_2$4668 and Mg$b$ are more suited to the high metallicity stellar population. The results suggest that principal component analysis method provides a more objective and informative alternative to diagnostics by individual spectral lines. ", + "introduction": "To understand how galaxies are formed and evolve we need to investigate their stellar populations and to derive their main parameters, such as metallicity and age. This study plays an important role in our understanding of galaxy properties (Kennicutt 1998; Maraston 1998; Padoan et al. 1997). During the last decade particularly using population synthesis method has performed determination of galaxy properties. It has been used extensively by many authors for all kinds of galaxies (Leitherer et al.\\ 1996). However, key parameters in interpretation of the observed properties are the metallicity and the age. The problem is the degeneracy of the effects from variations in age and in metallicity, even in the simplest unresolved stellar systems their effects are very difficult to separate (Vazdekis et al. 1997). It makes the determination of age and metallicity uncertain. To determine the age and metallicity more accurately, strong efforts have been devoted to select some spectral features that are more sensitive to age, and others are more sensitive to metallicity. In 1994, G. Worthey developed a very efficient method to investigate the age and metallicity effects in the integrated light of stellar populations. In his method, suppose an index varies by an amount $\\Delta I$, it can be explain $\\Delta I$ by either a pure metallicity variation or a pure age variation. The Z sensitivity parameter is the ratio of the percentage change in age to the percentage change in Z, with larger numbers indicating greater metallicity sensitivity (Worthey 1994). Using this method, some metallicity sensitive spectral indices are found, but only two age sensitive spectral indices are found. In this paper, we apply a different statistical technique, principal component analysis (PCA) to simple stellar population sample, to investigate on the reliability of the relations between spectral indices and age. We want to extract other spectral indices that can be used to determine the age of stellar population. The samples of simple stellar population (SSP) spectra come from Bressan et al.\\ (1996) and Bruzual \\& Charlot (1996). For each spectra, 21 spectral indices in the Lick/IDS system be measured first (Trager et al.\\ 1998), then we use PCA method to find some spectral indices that are more sensitive to the age. We leave the synthetic galaxy spectral indices with different star formation history and observational galaxy spectral indices to a future paper. However we already notice an encouraging good resemblance between the results of PCA and those one gets from observation (Worthey 1994). The organization of the rest of this paper is as follows. In section~2 we present PCA method. In section~3 we describe the definitions of spectral indices, the synthetic models we have used. In section~4 we analyze a sample of SSP spectra with different ages, find the first principal components (PCs), discuss the significance of different spectral indices to determine the age, and how different age sampling affects the results. In section~5, we illustrate that the uncertainties will not affect our results, and compare our results with previously published values. The conclusions are summarized in section~6. ", + "conclusions": "In order to select spectral indices that more sensitive to age than metallicity, PCA method has been applied to the Lick indices in the population model of SIL-SSP and BC-SSP. It is the first time this method is used for selecting the age sensitive spectral indices. Below, we summarize our main conclusions. \\begin{enumerate} \\item Using PCA method, we find some spectral indices, such as G4300, H$\\beta$, C$_2$4668 and Mg$b$, which can be used to determine the age of stellar population (1--15 Gyr). These spectral indices will help us to determine the age of stellar population more accurate. \\item Important differences are found between the age sensitive indices at different metallicities. For example, in low metallicity C$_2$4668 and Mg$b$ are not fit to determine the age of stellar population, but can be used to determine the age for high metallicity stellar population. \\item To understand how sensitive our method is to different stellar population synthesis models and metallicities, we have compared the results obtained with SIL-SSP and BC-SSP with 5 kind of metallicities. We find that although the precise values of PCs are difference, the general trend is very similar. We can conclude that the uncertainties in the evolutionary population synthesis models and metallicities can be ignored. \\end{enumerate}" + }, + "0112/astro-ph0112266_arXiv.txt": { + "abstract": "{ An optimization program is used to re-adjust the initial conditions, in order to reproduce as closely as possible the predictions of a complete ephemeris by using simplified equations of motion in the numerical integration. The adjustment of the initial conditions is illustrated in the transition from the DE406 complete long-range ephemeris to a Newtonian model considering only the Sun and the four major planets. It is also used to best reproduce this same DE406 ephemeris, based on post-Newtonian equations for a system of mass points and including the Moon and asteroids, by using a Newtonian calculation corrected by the Schwarzschild effects of the Sun and restricted to the ten major bodies of the solar system. ", + "introduction": "When integrating the equations of motion in classical or relativistic celestial mechanics, one has to know the values of the parameters, in particular the masses and the initial conditions. The masses of the celestial bodies are obtained by combining different methods, including the analysis of the perturbations that they cause to the motion of other bodies, and the fitting of the observed trajectory of a spacecraft in close approach. In the construction of modern ephemerides, the initial conditions are obtained by least-squares fittings to large sets of observational data \\citep{new, fie}. It is important to realize that the optimal parameters, that lead to the minimum (least-squares) residual with respect to a set of observational data, depend on the precise model that is used. Thus, in the case that an analytical ephemeris is adjusted to a numerically-integrated one, the initial conditions taken from the numerical ephemeris have to be modified \\citep{les, moi}. If one aims at testing an alternative theory of gravity, one may expect that even the masses will have to be very slightly modified \\citep{arm}. A program for the adjustment of the masses and the initial conditions has been built by the author in the latter context. It has been tested by adjusting a purely Newtonian calculation to the numerically-integrated ephemeris DE403 \\citep{sta95}, that is based on general-relativistic equations of motion. The program is based on the Gauss algorithm for iterative minimization of the mean quadratic error, which needs to calculate the partial derivatives of the theoretical predictions with respect to the parameters. Several standard methods may be used to compute these derivatives. In this program, they are calculated by finite differences: this involves making loops on the numerical solution of the equations of motion. In the present paper, this optimization program is used to re-adjust the initial conditions in the transition from the complete \"long\" ephemeris DE406 \\citep{sta98} to simplified models that are more tractable for integrations over very long times. Thus, the input data for the adjustment are taken from the complete ephemeris, and two simplified models of the solar system will be considered: in Section 2, we present the adjustment of a model based on purely Newtonian equations, and limited to the Sun and the four major planets. In Section 3, the initial conditions are re-adjusted for a model that considers the nine major planets and accounts, moreover, for those general-relativistic effects that are caused by the Sun alone. For these two adjustments, the sets of the optimal initial conditions are provided for potential users, and the differences with the predictions of the reference ephemeris are illustrated. Moreover, the effect of extrapolating the calculation outside the fitting interval can be seen, because the fitting intervals for the different planets are all significantly smaller than the 60 centuries of the DE406 ephemeris. ", + "conclusions": "The adequate values of all parameters in the equations of motion depend on the precise form adopted for the latter equations. In particular, it is very important to re-adjust the initial conditions when one passes from a complete form of the equations of motion, as it is used to build reference ephemerides, to a simplified form, which may be more adequate for long-range calculations. Although it may be too time-consuming to optimize the initial conditions so as to get the smallest residual for the whole time-range to be considered, one can get \"reasonably extrapolable\" initial conditions, by optimizing them for a time interval that contains at least a few periods of even the slowest body considered. Provided one thus re-adjusts the initial conditions, it is possible to reproduce quite accurately the corrections of post-Newtonian general relativity, by adding the Schwarzschild corrections, that account for the post-Newtonian effects of the Sun alone. This does not significantly increase the computer time as compared with the purely Newtonian equations of motion. A very precise comparison between the complete and simplified ways of describing post-Newtonian effects in the solar system could be obtained only if the two models were identical in the other respects, which was not the case here. Including the masses in the solved-for parameters brought only marginal improvements in the models for which it was considered." + }, + "0112/hep-ph0112279_arXiv.txt": { + "abstract": "\\noindent We investigate the dependence of the nucleon-nucleon force in the deuteron system on the values of coupling strengths at high energy, which will in general depend on the geometry of extra dimensions. The stability of deuterium at all times after nucleosynthesis sets a bound on the time variation of the ratio of the QCD confinement scale to light quark masses. We find the relation between this ratio, which is exponentially sensitive to high-energy couplings, and fundamental parameters, in various classes of unified theory. Model-dependent effects in the Higgs and fermion mass sector may dominate even over the strong dependence of the QCD scale $\\Lambda$. The binding energy of the deuteron also has an important effect on nucleosynthesis: we estimate the resulting bounds on variation of couplings. ", + "introduction": "In many models of particle physics, the universe is assumed to have more than 4 dimensions. The extra dimensions are either compactified to such a small size that we cannot (currently) probe them experimentally \\cite{polchinski} or possess metrics with a nontrivial dependence on the transverse directions such that we can only detect the gravitational influence of our familiar 4 dimensions \\cite{randall}. Cosmological solutions of the field equations of these theories often involve time evolution of the higher dimensions; the value of the gauge couplings in the low energy limit of these theories is invariably a function of the size or shape of the higher dimensions. This can be extremely problematic as variation in the gauge couplings over cosmological time scales may destroy the successful predictions of primordial nucleosynthesis \\cite{olive}. Recent claims of a time-variation of the fine structure constant \\cite{webb} also motivate the study of theories with dynamically-determined (and thus potentially time-varying) couplings. In relation to nucleosynthesis, it is only clear what the effect of changing one coupling constant at a time has upon the light element abundance. It is possible that degeneracies in the effect on nucleosynthesis may arise when more than one gauge coupling changes at once: an overall change in all couplings might be acceptable at a level much greater than that permitted for any one on its own. In this situation it is not possible accurately to constrain models such as \\cite{brax} where the gauge couplings oscillate with a fractional change of the order of 10\\% in the matter dominated era. Such a large fractional change would not be acceptable in, for example, the electromagnetic fine structure constant $\\alpha$ at nucleosynthesis, if this were the only time-dependent coupling. This paper is an attempt to provide an additional constraint, which is independent of nucleosynthesis (but which may affect calculations of nucleosynthesis), which suffers as little as possible from the problem of relating nuclear forces to underlying theory, and which is sensitive to a well-defined combination of couplings. Deuterium is only produced during nucleosynthesis, as it is too weakly bound to survive in the regions of stars where nuclear processes take place. The fact that deuterium is still observed today means that variations in the gauge coupling strengths or other fundamental parameters are non-trivially constrained by the requirement that the deuteron be stable at all times after nucleosynthesis. The fact that the deuteron is so weakly bound also makes it more sensitive to variations in the internuclear force. The strong running of $\\alpha_3$ at low energies means that a change in the coupling strength at high energy is manifested in a change in the strong coupling scale $\\Lambda_{\\rm QCD}$ (henceforth denoted by $\\Lambda$), by the usual dimensional transmutation arguments. Changes in $\\Lambda$ in turn lead to changes in the internuclear force. As recently pointed out by Langacker {\\em et al.}\\/\\ \\cite{LSS}, one also expects changes in quark masses and in the Higgs v.e.v.\\ $v$ if gauge couplings are unified at some scale. Any viable unified theory should accommodate (if not predict) a mechanism for electroweak symmetry-breaking, which may well depend sensitively on SUSY-breaking masses, and a mechanism of generating small Yukawa couplings, all of which may have a dependence on the unified coupling. Moreover, $\\Lambda$ is sensitive to all coloured particle masses through threshold effects of RG running. These effects introduce a large measure of model-dependence, since the correct theories of SUSY-breaking, electroweak symmetry-breaking and fermion masses are unknown. One may choose for simplicity to set to zero unknown effects in the electroweak, SUSY and Yukawa sectors \\cite{CalmetF}, but this runs the risk of neglecting terms which are of equal size or larger than the terms kept in the analysis. One might also consider ``less unified'' models, with more than one dynamically-determined fundamental coupling. The heterotic string dilaton $S$ and volume moduli $T$ provide a basic example, where gauge couplings and renormalisable Yukawa couplings (for canonically-normalised fields) have a universal dependence on $S$, but may be differently affected by changing $T$. The greater the number of independent quantities considered as time-dependent, the less predictive the theory becomes and the less meaningful are any constraints. Here we restrict ourselves to estimating the dependence of low-energy quantities in a somewhat idealised framework with a single dynamical unified coupling, corresponding to the v.e.v.\\ of a dilaton-like field. In the first part of this paper we calculate the deuteron binding energy by considering meson exchange forces, expressing the relevant parameters as a functions of $\\Lambda$ and the light quark masses $m_u$, $m_d$. The result is rather simple: we find that the deuteron is stable as long as the ratio $m_q/\\Lambda$ is greater than a certain value, where $m_q=m_u+m_d$. We perform similar calculations for the dineutron and diproton systems in the same isospin multiplet and investigate the criterion for their stability. Then we relate $\\Lambda$ and the quark masses to the QCD coupling strength and running masses at high energy using renormalisation group (RG) evolution. We take two cases, supersymmetric models with RG running up to the GUT scale (similar results will be obtained in the case of power-law unification in large extra dimensions), and nonsupersymmetric low-scale models with RG running up to a scale of a few TeV. The main result of this section is the exponential dependence of $\\Lambda$ on the perturbative strong coupling $\\alpha_3$ at high energy. We also find how the Higgs v.e.v.\\ and SUSY-breaking masses influence the low-energy parameters. Finally we consider how the bounds deduced from the two-nucleon system apply to various types of high-energy model, in which $\\alpha_3$ and the quark masses depend on model parameters (in particular the sizes of extra dimensions) which may be time-dependent. Thus, bounds on the possible cosmological evolution of such models since nucleosynthesis can be obtained. The constraints from the dinucleon system will in general apply to a different combination of theory parameters from those arising from nucleosynthesis --- taking the two together bounds the variation of fundamental parameters in two directions. (In addition, there are many other observational bounds applying at various much later epochs, discussed for example in \\cite{CalmetF}, including some from direct laboratory measurement.) We also point out for the first time and estimate the effect of changing the deuteron binding energy on the process of helium formation at nucleosynthesis, which may give rise to stronger bounds. \\subsection{Relation to recent work} The relation between a time-varying fine structure constant and other observables in particle physics was also treated in \\cite{BanksDD,CalmetF,LSS}. In \\cite{BanksDD} the resulting variation in the vacuum energy $V_0$ was estimated from general principles of QFT and found to be enormously larger than the cosmological bounds; hence, the authors concluded that a large number of implausibly accurate fine-tunings would be necessary for a time variation of the size that has been claimed to be consistent with field theory and cosmology. Such an argument is rather weak since it assumes, crucially, that the cosmological constant problem is ``solved'' at the present time by cancellations of different field theory contributions, all of which are many orders of magnitude larger than the measured value of $V_0$. This can hardly be a solid starting-point from which to set theoretical limits. The alternative considered in \\cite{BanksDD} was a self-tuning mechanism which protects the four-dimensional spacetime were we live by dynamically ``absorbing'' the vacuum energy into the curvature of one or more extra dimensions. Such a mechanism probably disallows inflation, but it is by no means clear that it also rules out spontaneous symmetry-breaking as claimed (the argument being that anything that ``cancels'' the vacuum energy also removes the source for a scalar to roll to the minimum of its potential). The existence of a dynamical time-scale for the self-tuning, the role of thermal effects in creating an effective potential, and the possibility that the extra-dimensional model may break four-dimensional Lorentz invariance are possible ways out of the $D=4$ field theory argument that $V(\\phi)$ must vanish at all times (see \\cite{Grojean} for related discussions). We prefer to take the majority point of view on the cosmological constant, {\\em i.e.} that it is our greatest theoretical area of ignorance and that no credible way to explain its smallness currently exists, therefore very little can be deduced from it, and certainly nothing related to quantum effects in the theory (which are the main difficulty in accommodating varying alpha). If we want to retain the semiclassical picture of matter coupled to gravity, the only sensible interpretation is a very light scalar, which interestingly would have about the same mass as quintessence; due to the coupling to electromagnetism, such a scalar would mediate composition-dependent forces which might have experimental signatures \\cite{DvaliZ}. But the vanishingly small scalar mass (even if experimentally confirmed) would remain a mystery, in the absence of an underlying theory which would explain why spacetime curvature was apparently so insensitive to a cosmologically-evolving field theory. Calmet and Fritszch \\cite{CalmetF} calculate some of the consequences for low-energy physics of changing $\\alpha$, within a GUT-like theory which constrains the SM gauge couplings to be equal at a particular energy scale (or at least to satisfy some fixed relation). They consider exclusively the effects on the strong interactions, with the unstated assumption that the mechanisms of electroweak symmetry-breaking, supersymmetry-breaking and fermion mass generation are held constant despite varying the unified coupling, hence that the quark and lepton masses, as well as $W$ and $Z$ masses, remain unchanged. Our calculation of the dependence of $\\Lambda$ and $M_N$ on $\\alpha_3(\\mu>m_t)$ is essentially identical to theirs, except for including the full dependence on thresholds. As noted above, such a study can only be a first step since the electroweak mass scales may also depend sensitively on the unified coupling (for example in models of hidden sector SUSY-breaking, radiative electroweak symmetry-breaking and anomalous U$(1)$ flavour models). Unfortunately, since we do not know the correct theories of electroweak and supersymmetry-breaking, let alone that of fermion masses, correlations between the time-variation of different low-energy observables (such as $\\alpha$ and the proton magnetic moment $\\mu_p$ or the mass ratio $m_p/m_e$) are necessarily model-dependent, unless one can find a combination insensitive to, say, the electroweak sector. We take the converse approach and quote in a hopefully less model-dependent way the bounds on high-energy parameters deriving from the low-energy system that we are studying. In particular one cannot claim yet (as in \\cite{CalmetF}) that the current data on the time variation of $\\alpha$ and other quantities are inconsistent with unification; but given an observational bound one can derive bounds on time-varying fundamental parameters, in whatever model one is interested in. With precise measurements of at least two different quantities at any particular epoch, of which at least one shows a nonzero variation, one {\\em can}\\/ rule out classes of unified theory that predict the wrong relation between (variations in the) different quantities. Time-varying couplings in principle are a new way of doing phenomenology, which allow us to test relations between the {\\em derivatives}\\/ of different quantities rather than their static values. This paper presents an additional bound which applies at all times after nucleosynthesis, thus for models in which time variation was more rapid in the early Universe (such as \\cite{brax}) it is likely quite restrictive. Langacker {\\em et al.}\\/\\ \\cite{LSS} have a somewhat similar approach to the ``high-energy'' aspect of the problem, parameterising the effects of variations in the unified coupling strength on the electroweak and Yukawa sectors by unknown, model-dependent constants of proportionality which depend on the particular model used. By looking at other precision astronomical measurements besides \\cite{webb}, they find that these constants must satisfy a rather precise relation, for the model to be consistent with observation. ", + "conclusions": "We have shown in this paper that a 3\\% increase in the QCD coupling constant $\\alpha_3$ at the GUT scale would result in the deuteron becoming unbound. The deuteron binding depends only on nuclear forces, so this conclusion cannot be escaped by considering the variation of more than one gauge coupling at once. Only negative variations at the level of 10\\% could bind the di-proton and the di-neutron. We have developed formalisms which enable one to calculate the variation in low energy parameters as a function of the variation of gauge and Yukawa couplings in the underlying theory: in many models one expects the variation in the electroweak and SUSY-breaking sectors to be the dominant effect at low energies, but the model-dependent nature of such effects means that no firm conclusions can yet be reached. We have also considered the effect of variations in the binding energy of the deuteron on the time at which helium formation occurs, and consequently on the helium abundance. This effect is complementary to the other effects on nucleosynthesis due to variation of gauge couplings, but on its own it constrains variation in $\\alpha_3$ at nucleosynthesis to within 0.25 \\%." + }, + "0112/astro-ph0112002.txt": { + "abstract": "The $\\approx$~1~Ms {\\it Chandra} Deep Field North observation is used to study the extended X-ray sources in the region surrounding the Hubble Deep Field North (HDF-N), yielding the most sensitive probe of extended X-ray emission at cosmological distances to date. A total of six such sources are detected, the majority of which align with small numbers of optically bright galaxies. Their angular sizes, band ratios, and X-ray luminosities --- assuming they lie at the same distances as the galaxies coincident with the X-ray emission --- are generally consistent with the properties found for nearby groups of galaxies. One source is notably different and is likely to be a poor-to-moderate X-ray cluster at high redshift (i.e., $z\\ga0.7$). This source has a large angular extent, a double-peaked X-ray morphology, and an overdensity of unusual objects [Very Red Objects, optically faint ($I\\ge24$) radio and X-ray sources]. Another of the six is coincident with several $z\\approx1.01$ galaxies located within the HDF-N itself, including the FR~I radio galaxy VLA~J123644.3+621133, and is likely to be a group or poor cluster of galaxies at that redshift. We are also able to place strong constraints on the optically detected cluster of galaxies ClG~1236$+$6215 at $z=0.85$ and the wide-angle-tail radio galaxy VLA~J123725.7$+$621128 at $z\\sim1$--2; both sources are expected to have considerable associated diffuse X-ray emission, and yet they have rest-frame 0.5--2.0~keV X-ray luminosities of $\\la3\\times10^{42}$\\xlum and $\\la(3$--$15)\\times10^{42}$\\xlum, respectively. The environments of both sources are either likely to have a significant deficit of hot intra-cluster gas compared to local clusters of galaxies, or they are X-ray groups. We find the surface density of extended X-ray sources in this observation to be 167$^{+97}_{-67}$ deg$^{-2}$ at a limiting soft-band flux of $\\approx3\\times10^{-16}$\\xflux. No evolution in the X-ray luminosity function of clusters is needed to explain this value. ", + "introduction": "The Hubble Deep Field North \\citep[HDF-N;][]{Williams1996,Ferguson2000} was chosen as the location of a deep {\\it Hubble Space Telescope} survey because it contained no known bright sources at radio, infrared, optical, or X-ray wavelengths and no known nearby ($z<0.3$) clusters of galaxies. This effort was conceived to advance the study of galaxy evolution to high redshifts, but it has since initiated one of the most intensive, multi-wavelength investigations on the sky, influencing a wide range of astronomical topics (e.g., Ferguson et al. 2000 and references therein) and yielding one of the most comprehensive data sets publicly available (e.g., deep imaging at nearly all astronomically accessible wavelengths and more than 700 spectroscopic redshifts within a $\\sim4\\arcmin$ radius of the HDF-N).\\footnote{See the Hubble Deep Field Clearinghouse; http://www.stsci.edu/ftp/science/hdf/clearinghouse/clearinghouse.html.} Recently, the {\\it Chandra X-ray Observatory} \\citep[hereafter {\\it Chandra};][]{Weisskopf2000} completed an $\\approx$~1~Ms survey of the HDF-N and its environs, providing an extremely sensitive view of the X-ray Universe. The {\\it Chandra} Deep Field North Survey \\citep[CDF-N;][hereafter Paper~V]{Brandt2001b} covers an area approximately $18\\arcmin \\times 22\\arcmin$ in size and reaches 0.5--2.0~keV (soft) and 2.0--8.0~keV (hard) flux limits of $\\approx 3\\times 10^{-17}$\\xflux\\ and $\\approx 2\\times 10^{-16}$\\xflux\\ (5.5$\\sigma$) near the aimpoint for point sources. In addition to resolving most of the X-ray background into individual point sources \\citep[e.g., Paper V,][]{Cowie2002}, the observation allows the detection of even relatively poor clusters and groups of galaxies to significant redshifts ($z\\sim1$). The detection of extended X-ray emission from hot gas provides compelling evidence that apparent optical clusters or groups are truly gravitationally bound. Since the properties of clusters and groups are intimately dependent upon cosmological parameters and the history of structure formation, X-ray emission offers a useful probe of hierarchical structure. There is already substantial observational evidence that sources in the vicinity of the HDF-N tend to cluster at certain redshifts in ``walls'' or ``filaments,'' and that a few of these redshift peaks can be broken up spatially into apparent groups of galaxies \\citep[e.g.,][]{Cohen2000, Dawson2001}. The CDF-N has the potential to determine whether the gravitational potential wells of these apparent groups are deep enough to harbor large amounts of hot gas and dark matter. Additionally, deep radio imaging of this region has revealed two highly extended radio sources \\citep{Richards1998, Muxlow1999, Snellen2001}; one is the Fanaroff-Riley~I \\citep[FR~I;][]{Fanaroff1974} radio galaxy VLA~J123644.3+621133 located within the HDF-N itself, and the other is the wide-angle-tail (WAT) source VLA~J123725.7$+$621128. These two radio \\end{multicols} \\vbox{ \\vspace{17.3cm} %\\centerline{ %\\hglue-1.2in{\\includegraphics[width=21cm]{fig1.color.eps}} %} \\vspace{0.0cm} \\figcaption[fig1.color.eps] {Adaptively smoothed ``true color'' X-ray image with red, green, and blue representing 0.5--2.0~keV, 2.0--4.0~keV, and 4.0--8.0~keV emission, respectively. Prior to combination, each color image was smoothed to 2.5$\\sigma$ with an adaptive kernel algorithm \\citep{Ebeling2001} to enhance diffuse features. The white boxes denote the six extended X-ray sources detected (see $\\S$\\ref{notes}), while the two dashed circles indicate cluster candidates found at other wavelengths that show no significant extended X-ray emission (see $\\S$\\ref{nondetections}). In order to show clearly the extended sources, it was necessary to saturate bright X-ray point sources. This tends to make all of the brighter point sources appear yellow in the image, regardless of whether their spectral shapes are particularly hard or soft.\\label{fig:xrayfig}} } \\begin{multicols}{2} \\noindent sources are notable because FR~Is, and in particular WATs, are known to reside predominantly in or near rich clusters of galaxies; extended X-ray emission associated with these particular radio sources has yet to be detected. Arguably the most definitive evidence for clustering near the HDF-N, however, comes from \\citet{Dawson2001} who recently reported the discovery of the $z=0.85$ cluster of galaxies ClG 1236$+$6215 slightly north of the HDF-N. Dawson et al. predict this cluster should have a bolometric X-ray luminosity in excess of $10^{44}$\\xlum, which, for an extended object of radius 30$\\arcsec$ at a redshift of $z=0.85$, should be nearly a factor of 15 above the detection threshold of the CDF-N. Although a 21~ks {\\it ROSAT} High Resolution Imager (HRI) observation of this region detected no extended X-ray objects, the resulting soft-band flux threshold of $\\sim$2$\\times10^{-14}$\\xflux\\ is not particularly constraining; for instance, the HRI observation would have missed a typical $10^{43}$\\xlum\\ X-ray cluster at $z\\ga0.4$ or even a typical $10^{44}$\\xlum\\ X-ray cluster at $z\\ga0.95$. The 1~Ms {\\it Chandra} observation allows us to push these constraints much lower. The deepest soft-band X-ray observations prior to {\\it Chandra} were those of the {\\it ROSAT} Ultra Deep Survey toward the Lockman Hole \\citep[hereafter UDS; e.g.,][]{Lehmann2001}, which detected ten extended objects over a $\\sim$$30\\arcmin$ field of view. A few of these sources are classified as clusters \\citep[including the double-peaked lensing cluster RX~J105343$+$5735;][]{Hasinger1998}, but the majority appear to be groups; all are thought to lie at redshifts of $z\\sim0.2$--1.0. Results from other deep surveys with {\\it ROSAT} \\citep[e.g.,][]{McHardy1998, Zamorani1999} were comparable. While the CDF-N covers only $\\approx$~$\\onequarter$ the area of the UDS, it probes the X-ray sky a factor of $\\approx~7$ times deeper for extended objects in the soft band, making the discovery of fainter and more distant objects possible. Here we describe the nature of the extended X-ray sources detected within the extremely deep CDF-N observation. In $\\S$\\ref{data} and $\\S$\\ref{optical}, we briefly outline our reduction and detection techniques for the X-ray and optical observations, respectively. Descriptions of individual sources are presented in $\\S$\\ref{notes}. Finally our findings are discussed and summarized in $\\S$\\ref{discuss}. Throughout this paper, we adopt $H_{0}=65$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_{\\rm M}=\\onethird$, and $\\Omega_{\\Lambda}=\\twothirds$. The Galactic column~density toward the CDF-N is $(1.6\\pm0.4)\\times10^{20}$ cm$^{-2}$ \\citep{Stark1992}. Coordinates are for the J2000 epoch. ", + "conclusions": "\\label{discuss} \\subsection{Basic Nature of the Extended X-ray Sources}\\label{basic} The general X-ray and optical characteristics of the extended sources in the CDF-N (i.e., their soft X-ray luminosities, apparent X-ray sizes, and weak optical clustering) are most comparable to those of nearby groups of galaxies \\citep[e.g.,][]{Mulchaey1996, Ponman1996, Helsdon2000}. The only exception appears to be source~2, which has a larger angular size and contains an overdensity of unusual objects; it is likely to be a poor-to-moderate X-ray cluster at high redshift (i.e., $z\\ga0.7$). This CDF-N observation has also allowed us to place strong constraints on two potential extended X-ray emitting systems, ClG 1236$+$6215 and VLA~J123725.7$+$621128. Both systems lie at high redshift and, from their undetected status, may still be in the early stages of dynamical evolution, having not yet formed a central concentration of hot gas and dark matter massive enough to produce detectable X-ray emission \\citep[e.g.,][]{Blanton2001, Donahue2001}. For the two sources with enough counts to perform detailed spectral analysis (sources~2 and 6), we have modeled their X-ray spectra and estimated thermal plasma temperatures. Comparing the most likely rest-frame X-ray luminosities and temperatures of these two sources with the well-established $L_{\\rm X}-kT$ relation for clusters and groups \\citep[e.g.,][]{Allen1998, Xue2000}, we find that the X-ray temperature of source~2 is consistent with its bolometric luminosity if the X-ray source lies at high redshift (i.e., $z\\ga0.7$), while the X-ray temperature of source~6 is clearly too high to be consistent with its bolometric X-ray luminosity \\citep[compare to Figure~1 of][]{Xue2000}. Whether the high temperature of source~6 is due point-source contamination or non-gravitational heating mechanisms such as star formation from individual group members or shock heating of the infalling gas \\citep[e.g.,][]{Metzler1994, Ponman1996, Cavaliere1997} cannot be resolved with our current X-ray data. Some additional knowledge about the state of the X-ray emitting gas in source~2 can be gained from its observed X-ray morphology. The source appears to be irregular and double-peaked, suggesting a young merger. Unlike source~6, this double-peaked source appears to retain its extended, bimodal structure even in the minimally smoothed X-ray images (e.g., similar to the images made for source~6 in $\\S$\\ref{detections}) and is likely to be real.\\footnote{Although we cannot conclusively exclude that possibility of some faint, point-source contamination.} Recent simulations of offset merging clusters \\citep{Ricker2001} suggest that during the evolution of merging clusters, there is a short-lived phase of increased luminosity and temperature. Furthermore, for large impact parameters, the morphological structure of the merger remnant becomes bimodal. While this phase is likely to last at most only a few Gyrs, the strong enhancement in luminosity may help to offset their low observational occurrence. Perhaps the bimodal structure we see in this source~2, as well as in that of the Lockman Hole source RX~J105343$+$5735, are a natural consequence of this presumably common, albeit short-lived, stage in the formation of clusters and groups. \\subsection{Number Density of Extended X-ray Sources}\\label{density} The six detected sources are all found in the ``high-exposure'' area ($\\approx$~130 arcmin$^{2}$; see $\\S$\\ref{data}) with exposure times above 800~ks, implying an extended-source surface density of 167$^{+97}_{-67}$ deg$^{-2}$ (1$\\sigma$) at a limiting soft-band flux of $\\approx3\\times10^{-16}$\\xflux. \\vbox{ \\vspace{-1.65cm} \\centerline{ \\hglue0.1cm{\\includegraphics[width=9cm]{fig6.eps}} } \\vspace{-1.5cm} \\figcaption[fig6.eps]{Cumulative number counts as a function of soft-band X-ray flux. Plotted are the extended sources from this sample as well as those of \\citet{McHardy1998}, \\citet{Vikhlinin1998}, \\citet{Zamorani1999}, and \\citet{Lehmann2001}. The solid and dashed lines show two models of the number counts derived from integrating the XLF of \\citet{Ebeling1997} above $L_{\\rm 0.5-2.0~keV}=1.3\\times10^{42}$\\xlum\\ and $1.2\\times10^{41}$\\xlum, respectively (see $\\S$\\ref{density} for details).\\label{fig:logN-logS}} } \\vspace{0.4cm} \\noindent Figure~\\ref{fig:logN-logS} shows the cumulative soft-band number counts for \\hbox{CDF-N} extended sources (open circles). Also shown are the number counts from the {\\it ROSAT} extended-source samples of \\citet{McHardy1998}, \\citet{Vikhlinin1998}, \\citet{Zamorani1999}, and \\citet{Lehmann2001}. For comparison, we calculated the expected number density of X-ray clusters using the Schechter expression from the bolometric parameterization of the local cluster X-ray luminosity function (XLF) by \\citet[][see their Table~1]{Ebeling1997}. Two calculations were made. In both cases, we assumed no evolution in the $\\Omega_{\\rm M}=\\onethird$, $\\Omega_{\\Lambda}=\\twothirds$ cosmology, and the integration was made over the range $z=0.015$--1.2. We assigned a plasma temperature to each cluster assuming its bolometric luminosity follows the $L_{\\rm X}-kT$ relations of \\citet{Xue2000}. The solid line in Figure~\\ref{fig:logN-logS} denotes the expected number counts for sources with $L_{\\rm X}>2.5\\times10^{42}$\\xlum\\ (bolometric, or equivalently $1.3\\times10^{42}$\\xlum\\ in the 0.5--2.0~keV band for $kT=1$~keV). Above this luminosity, the Ebeling et al. XLF is supported by observational data \\citep[see also][]{Rosati1998}. Only two of the six CDF-N extended sources (or 55 sources deg$^{-2}$) are likely to lie above this luminosity limit; this number density is consistent with the predicted model. The dashed line in Figure~\\ref{fig:logN-logS} shows an extrapolation of the Ebeling et al. XLF to include objects with $L_{\\rm X}>2.0\\times10^{42}$\\xlum\\ (bolometric, or $1.2\\times10^{41}$\\xlum\\ in the 0.5--2.0~keV band for $kT=0.5$~keV). Note that systematic studies of extended X-ray sources with $L_{\\rm X}\\la5.0\\times10^{42}$\\xlum\\ have been hindered by their faint X-ray fluxes even locally, so this XLF extrapolation, while plausible, is somewhat uncertain. Even so, these faint sources appear to be consistent with the predicted model. Thus the CDF-N sources appear to be consistent to within error with no evolution in the XLF. We caution, however, that our results may be affected by ``cosmic variance'' due to the small field of view of this observation. Observations of larger-area fields with {\\it Chandra} and {\\it XMM-Newton} are needed to confirm this result. Most of the extended sources in the CDF-N are likely to be groups, and, as such, are thought to be affected more by energy and momentum feedback from the stellar winds and supernovae following star formation than by cosmological parameters \\citep[e.g.,][]{Cavaliere2000}. In fact, such pre-heating phase is currently the best method for reconciling the different $L_{\\rm X}-kT$ relations of groups and clusters. The lack of evolution in the extended source number counts suggests that this pre-heating phase may occur at redshifts higher than we tentatively observe here. These findings also show that extended sources with fluxes $\\la10^{-15}$\\xflux\\ are not likely to contribute significantly to the 0.5-2.0~keV cosmic X-ray background \\citep[$\\la0.7$\\% using the results of][]{Cowie2002}. By comparison, extended sources with fluxes $\\ga10^{-14}$\\xflux\\ are thought to contribute $\\sim$10\\% to the cosmic X-ray background \\citep[e.g.,][]{Oukbir1997, Rosati1998}. The fact that we do not detect extremely luminous sources is understandable given their low surface density and that the HDF-N was chosen to avoid such objects. \\subsection{Coincidence of X-ray Emission with Other Clustering Indicators}\\label{peaks} X-ray emission is thought to be an efficient method for tracing intermediate-density structure within the ``cosmic web'' of large-scale structure \\citep[e.g.,][]{Tully1987, Bond1996, Pildis1996, Cen1999, Pierre2000}. The extensive multi-wavelength observations within the HDF-N and surrounding regions afford us an opportunity to compare the clustering seen at X-ray wavelengths with that seen using optical-to-near-IR and radio methods. We find significant overdensities of optical galaxies around most of the CDF-N extended X-ray sources from Monte Carlo simulations, but see little evidence for any strong red sequence of early type galaxies and, perhaps surprisingly, a stark contrast between the X-ray and optical appearances of detected clusters in Figure~\\ref{fig:overlay}. Unlike most local clusters or groups, very few of the extended X-ray sources appear to be associated with obvious optical clusters or groups. This outcome is quite understandable, given the dramatic increase in the field galaxy number density at these magnitudes, but highlights a potential deficiency in relying on optical methods alone. Considering the two sources with secure redshifts, source~3 falls on one of the redshift peaks noted by \\citet{Cohen2000} and \\citet{Dawson2001}, and source~6 does not. This suggests that at least some low-level clustering occurs within these observed ``filaments'' of galaxies. The redshifts of the four other extended X-ray sources are still unknown and may possibly coincide with known overdensities. If the coincidence of X-ray groups and clusters with known ``filaments'' in the CDF-N region is anything like that in the {\\it ROSAT} North Ecliptic Pole Survey, we should expect at least 25\\% (or $\\sim$2) of the extended X-ray sources to lie within such large-scale structures \\citep[e.g.,][]{Burg1992, Brinkmann1999, Mullis2001}. Alternatively, we find no evidence for X-ray emission associated with the optically detected cluster ClG 1236$+$6215. At radio wavelengths, two FR~I radio sources are known within the CDF-N region. These sources are predominantly thought to reside in or near rich clusters of galaxies. We find that one of these is radio objects is coincident with an extended X-ray source near the threshold of our current {\\it Chandra} observation, while the other remains undetected (although X-ray and optical constraints cannot rule out a $z>1$ cluster). Clearly the spatial coverage of the CDF-N is too small to permit firm conclusions about selection biases between these different cluster indicators. However, the analysis presented here does suggest that relying on any single technique may fail to detect some (and possibly many) potential clusters and groups. \\subsection{Future Work}\\label{future} We have provided here an extremely deep view of the faint extended X-ray source population in the CDF-N. Several details about these objects are revealed, particularly for the brightest two sources, one of which appears to be bimodal and is perhaps undergoing a merger. While the six CDF-N objects generally confirm our understanding of less distant clusters and groups, they also provide a look into the past, when such objects were just beginning to form. Tighter constraints on these sources, both in terms of accurate distance determinations and deeper optical imaging, should improve our understanding of the filamentary structures in the vicinity of the HDF-N and (perhaps) cosmological models. Secure optical counterparts for many of the CDF-N extended sources are clearly the most important piece of information currently missing from our picture of these distant systems. Obtaining further constraints on the X-ray properties of these sources (e.g., temperatures, abundances, surface density profiles) will be difficult, given the large amount of observing time needed simply to detect them. As observations of the CDF-N continue with both {\\it Chandra} and {\\it XMM-Newton}, it may be possible to search for temperature variations between the two peaks in the tentative high redshift system. This would be a valuable diagnostic for evaluating the nature of this extended X-ray source (e.g., is the variation consistent with a merger shocks or cold cluster cores moving through low density, shock-heated intra-cluster gas)." + }, + "0112/astro-ph0112521_arXiv.txt": { + "abstract": "\\noindent $``$Analyzing a spectrum is exactly like doing a crossword puzzle, but when you get through with it, you call the answer research.\" {\\it -- Henry Norris Russell.} \\bigskip The absorption lines observed in quasar spectra have given us a detailed picture of the intergalactic medium and the metal abundance and kinematics of high redshift galaxies. In this review, we present an introduction to the field, starting with the techniques used for interpreting absorption line spectra. We then survey the observational and theoretical development of our understanding of the Lyman $\\alpha$ forest, the metal absorbers, and the damped Ly$\\alpha$ absorbers. We conclude with a discussion of some of the remaining outstanding issues, and prospects for the future. ", + "introduction": "Absorption lines were seen in the earliest photographic spectra of quasars in the mid-1960's (Sandage 1965, Gunn $\\&$ Peterson 1965, Kinman 1966, Burbidge, Lynds $\\&$ Burbidge 1966, Burbidge 1967). By 1969, the concensus was that quasar redshifts are cosmological, and that many of the absorption lines in quasar spectra originate in intervening galaxies (see Bahcall \\& Salpeter 1965, Bahcall \\& Spitzer 1969). Subsequently, W. L. W. Sargent, Peter J. Young and collaborators wrote a series of papers presenting the first comprehensive studies of QSO absorption lines, using high quality spectra obtained with Boksenberg's photon counter, the IPCS, at the Double Spectrograph on the Palomar Hale telescope (Young et al. 1979; Sargent et al. 1979; Sargent, Young, Boksenberg \\& Tytler 1980; Young, Sargent \\& Boksenberg 1982ab; Sargent, Young \\& Schneider 1982). A classification scheme was developed, summarized in an article in the Annual Review of Astronomy and Astrophysics, by Weymann, Carswell \\& Smith (1981). First, the so-called $``$metal-line absorbers\" with redshifts, $z_{abs}$, much less than the quasar emission line redshift, $z_{em}$, were attributed to interstellar gas in intervening galaxies. Systems which were detected in Ly$\\alpha$ and occasionally the Lyman series of hydrogen, were dubbed the Ly$\\alpha$ forest. These showed no clustering like galaxies, nor detectable metals. They were thought to be primordial, pre-galactic, inter-galactic gas, pressure-confined by a hypothesized inter-galactic medium, and ionized by the integrated UV radiation of quasars themselves (Sargent, Young, Boksenberg \\& Tytler 1980). Finally, some quasars showed broad, P-Cygni like troughs of absorption at the emission line redshift; these were interpreted as radiatively driven winds associated with the quasar central engine (Turnshek 1984 and references therein). Quasar absorbers are useful probes of the high redshift universe for many reasons. Atoms have a rich ultraviolet absorption spectrum; typically a single spectrum will contain lines from many elements in a range of ionization states. Lines like Ly$\\alpha$ are extremely sensitive probes of small amounts of gas. Absorption spectra are relatively easy to interpret, since the lines arise from atoms in the ground state, and probe a pencil beam through the volume of gas. By contrast, emission lines are a complicated integral of density and other physical conditions over the emitting volume. Until recently, quasar absorbers contained the only information about $``$normal\", that is, non-active galaxies at high redshift. They still are the best way to study the detailed kinematics and metal abundances of high redshift galaxies. Other reviews of the subject can be found in the conference proceedings edited by Meylan (1995) and Petitjean $\\&$ Charlot (1997). The Ly$\\alpha$ forest was reviewed by Rauch (1998) and damped Ly$\\alpha$ absorber abundunces by Lauroesch et al (1996). ", + "conclusions": "" + }, + "0112/astro-ph0112467_arXiv.txt": { + "abstract": "The aim of this paper is to describe new statistical methods for determination of the correlations among and distributions of physical parameters from a multivariate data with general and arbitrary truncations and selection biases. These methods, developed in collaboration with B. Efron of Department of Statistics at Stanford, can be used for analysis of combined data from many surveys with different and varied observational selection criteria. For clarity I will use the luminosity function of AGNs and its evolution to demonstrate the methods. I will first describe the general features of data truncation and present a brief review of past methods of analysis. Then I will describe the new methods and results from simulations testing their accuracy. Finally I will present the results from application of the methods to a sample of quasars. ", + "introduction": "One of the important ways of testing the models of AGN, or any other astrophysical source, is through the investigations of the distributions, ranges, and more importantly the correlations among, the relevant physical characteristics, such as luminosity, spectrum, redshifts or distances. A reliable determination of these features requires large samples. As evident from papers presented in this proceedings the samples are becoming larger and larger. Combining the samples, however, is a very challenging task, because different samples are obtained by different instruments and techniques, so that they suffer from different and varied selection biases and data truncations. Overcoming these biases requires care. The primary goal of this paper is to describe some of the relatively new methods we have developed at Stanford over the past decade (Efron \\& Petrosian 1992, 1999). These methods are very general and are applicable to any data with well defined but arbitrary truncations. We have applied these to various astrophysical data such as solar flares (Lee, Petrosian, \\& McTiernan 1993, 1995), gamma-ray burst (see e.g. Lloyd, Petrosian, \\& Mallozzi 2000) and quasars (Meloney \\& Petrosian 1999). Instead of using abstract mathematical symbols, the method will be demonstrated using the luminosity function of AGNs and its cosmological evolution, {\\it i.e.} its variation with redshift $z; \\Psi(L,z)$. Without loss of generality, we can write the luminosity function as \\beq\\label{lf} \\Psi(L,z) = \\rho (z) \\psi(L/g(z),\\alpha_i) / g(z), \\eeq where $\\rho(z)$ describes the co-moving density evolution and $g(z)$ (with $g(0) = 1$) describes the luminosity evolution of the population with $L_o = L/g(z)$ as the luminosity adjusted to its present epoch value; $\\psi (L_o, \\alpha_i)$ gives the local luminosity function. Here I explicitly include the shape parameters $\\alpha_i$, which could also vary with redshift. A surprising result has been the absence of evidence for a strong shape evolution. In this paper I ignore such effects and concentrate on the determination of the the density and luminosity evolution functions $\\rho(z)$ and $g(z)$. In the next section we describe various kind of truncations and in \\S 3 give a brief summary of some of the past methods used for this kind of analysis. The new methods and their accuracy are described in \\S 4 and a sample of results from application to AGNs are summarized in \\S 5. ", + "conclusions": "" + }, + "0112/astro-ph0112184_arXiv.txt": { + "abstract": "In this Paper we present the source catalog obtained from a 942 ks exposure of the \\chandra Deep Field South (CDFS), using the Advanced CCD Imaging Spectrometer (ACIS--I) on the \\chandra X--ray Observatory. Eleven individual pointings made between October 1999 and December 2000 were combined to generate the final image used for object detection. Catalog generation proceeded simultaneously using two different methods; a method of our own design using a modified version of the {\\tt SExtractor} algorithm, and a wavelet transform technique developed specifically for \\chandra observations. The detection threshold has been set in order to have less than 10 spurious sources, as assessed by extensive simulations. We subdivided the catalog into four sections. The primary list consists of objects common to the two detection methods. Two secondary lists contain sources which were detected by: 1) the {\\tt SExtractor} algorithm alone and 2) the wavelet technique alone. The fourth list consists of possible diffuse or extended sources. The flux limits at the aimpoint for the soft (0.5--2 keV) and hard (2--10 keV) bands are 5.5$\\times10^{-17}$ erg s$^{-1}$ cm$^{-2}$ and 4.5$\\times10^{-16}$ erg s$^{-1}$ cm$^{-2}$ respectively. The total number of sources is 346; out of them, 307 were detected in the 0.5--2 keV band, and 251 in the 2--10 keV band. We also present optical identifications for the catalogued sources. Our primary optical data is $R$ band imaging from VLT/FORS1 to a depth of $R \\sim 26.5$ (Vega). In regions of the field not covered by the VLT/FORS1 deep imaging, we use $R$--band data obtained with the Wide Field Imager (WFI) on the ESO--MPI 2.2m, as part of the ESO Imaging Survey (EIS), which covers the entire X--ray survey. We found that the FORS1/Chandra offsets are small, $\\sim 1\\arcsec$. Coordinate cross-correlation finds $85\\%$ of the \\chandra sources covered by FORS1 $R$ to have counterparts within the 3$\\sigma$ error box ($\\gtrsim1.5\\arcsec$ depending on off--axis angle and signal--to--noise). The unidentified fraction of sources, approximately $\\sim$ 10--15\\%, is close to the limit expected from the observed X--ray flux to $R$--band ratio distribution for the identified sample. ", + "introduction": "} The X--ray background (hereafter XRB) was first identified by \\citet{giacconi1962} and its resolution into discrete sources has since been a major goal of X--ray astronomy. Deep surveys by the major X--ray facilities UHURU \\citep{matilsky1973}, HEAO-1 A2 \\citep{piccinotti1982}, Einstein \\citep{giacconi1979}, ROSAT \\citep{hasinger1998}, ASCA \\citep{ueda1999} and BeppoSAX \\citep{giommi2000} have resolved an increasing fraction of the XRB in the 0.5--2 keV and in the 2--7 keV band and produced catalogs of sources which were then used for a wide variety of other astrophysical investigations. These surveys were primarily limited by effective area and source confusion at the faintest achievable fluxes. Unlike these previous missions, \\chandra \\citep{weiss2000} provides arcsecond resolution over a majority of the detector, in particular the Advanced CCD Imaging Spectrometer \\citep[ACIS--I, ][]{garmire1992, bautz1998}. Superior spatial resolution coupled with a substantial gain in light--gathering power over ROSAT has allowed our 942 ks exposure (hereafter 1Msec) of the CDFS (and the analogous exposure of the Hubble Deep Field North) to become the deepest X--ray exposure(s) ever taken, improving by factors of $\\sim20$ and $\\sim200$ the deepest ROSAT and ASCA surveys, respectively. Along with the \\chandra Deep Field North (CDFN) \\citep{hornschemeier2000,hornschemeier2001, brandt2001a,brandt2001b}, the CDFS provides a unique dataset in which to investigate both statistical and source-by-source properties of active galaxies over a large range of redshift and parameter space. Already, the CDFS has discovered both single interesting objects \\citep{norman2001} and statistical correlations (Giacconi et al. 2001; Tozzi et al. 2001, hereafter Paper I and II). The \\chandra Deep Field South is centered on $\\alpha$ = 03:32:28.0, $\\delta$=-27:48:30 (J2000), and was selected as having: (1) low Galactic neutral hydrogen column (N$_H$ $\\sim$ 8 $\\times$ 10$^{19}$ cm$^{-2}$); (2) no bright stars ($m_{\\textrm v} \\leq 14$) within 30$\\arcmin$ and (3) field accessibility from the new 8m class telescopes, namely VLT and Gemini-South. Note that recent higher--resolution HI maps (from the Parkes Multi--beam Survey; L. Staveley--Smith, private communication) confirm the HI hole in the CDFS, but do show some sub--structure (on $\\sim 10$ arcminute scale) near the \\chandra pointings, at the level of $10^{19.5}$ cm$^{-2}$. Here we present the catalog of the sources in the CDFS field found in the 1Msec exposure along with fluxes (or upper limits) in $R$, for their presumed optical counterparts. The first scientific results derived from the 1Msec observation are presented in Rosati et al. 2001 (Paper III). The Paper is structured as follows: in Section \\ref{data} we discuss the data, in Section \\ref{detection} we discuss extraction for point and extended sources and their photometry; in Section \\ref{opticalIDs} the optical identifications are presented. We conclude with a discussion of this dataset. The common catalog is given in Table~\\ref{common}. In Tables \\ref{sex}, \\ref{wav} and \\ref{extended} are listed sources detected only by SExtractor and wavelet techniques, and extended properties respectively. ", + "conclusions": "We briefly summarize the detection criteria adopted in this Paper. There are three basic steps: 1) a detection algorithm (either SExtractor or WAVDETECT) has been run on the full 0.5--7 keV image to generate a list of source candidates; 2) aperture photometry has been performed independently in the 0.5--2 keV image and 2--7 keV image at the position of each source candidate; 3) all those sources with S/N$>$2.1 either in the soft or in the hard image are considered as real sources and appear in our catalogs. This procedure may not reach the deepest sensitivity achievable with our observations, however, it is a good compromise achieving low limiting fluxes without introducing numerous fake sources. As assessed by extensive simulations described in \\citet{tozzi2001}, the fraction of fake sources to be expected in our sample is less than 10 (or $\\sim 3$\\%) with this method. The faintest sources in our samples have approximately 10 net counts. Given the exposure time and the adopted conversion factors the limiting fluxes achieved in the 0.5--2 keV and in the 2--10 keV band are $5.5\\times10^{-17}$ erg s$^{-1}$ cm$^{-2}$ and $4.5\\times10^{-16}$ erg s$^{-1}$ cm$^{-2}$, respectively. These limiting fluxes are about 20 and 200 times lower than that previous X--ray missions in the soft and hard bands respectively. The catalog of sources presented in this Paper allows us to investigate the nature of X--ray faint sources with unprecedented sensitivity, together with their optical properties. Most of the sources in our sample fall within the X--ray to optical flux ratio range typical of AGN as determined from the EMSS sources (Stocke et al. 1991). In addition a significant population of X--ray faint/optically bright sources is observed, with $f_x/f_R\\leq 0.1$, lower than those typical of AGNs. Most of these sources are detected in the soft band only (circles in Figure ~\\ref{fxfop}). This population of sources has been already partially identified as nearby, bright normal galaxies by \\citet{tozzi2001,hornschemeier2001,barger2001}. Only a few sources seem to have $f_{x}/f_{R}$ values higher than usual, which may be obscured AGN at very high redshift. This population of optically faint sources ($R>25$) is consistent with a mix of obscured AGN at $z=1-3$ and evolved, high--z galaxies (see Alexander et al. 2001; Cowie et al. 2001). The nature of these sources will be investigated in greater detail in following papers. The data presented here constitute (along with the \\chandra Deep Field North) the deepest X--ray exposure ever taken. As such, it is a unique dataset for current and future research. A proposal to re--observe this field with \\chandra will await detailed analysis of the current dataset. However, a SIRTF Legacy program will be observing this field with both MIPS and IRAC during their own deep survey. Also, a 500ks observation with XMM has been planned in the CDFS this year. The high XMM throughput, combined with the arcsec \\chandra resolution, will allow high quality X--ray spectra and secure optical identifications to be obtained for most of the X--ray sources in the CDFS. The intensive optical and infrared coverage is planned to continue and to be vigorously pursued at ESO. Deep VLA radio observations are being analyzed." + }, + "0112/hep-ph0112222_arXiv.txt": { + "abstract": "We compare the tau neutrino flux arising from the galaxy and the earth atmosphere for $10^{3} \\leq E/\\mbox{GeV} \\leq 10^{11}$. The intrinsic and oscillated tau neutrino fluxes from both sources are calculated. The intrinsic galactic $\\nu_{\\tau}$ flux ($E \\geq 10^{3}$ GeV) is calculated by considering the interactions of high-energy cosmic-rays with the matter present in our galaxy, whereas the oscillated galactic $\\nu_{\\tau}$ flux is coming from the oscillation of the galactic $\\nu_{\\mu}$ flux. For the intrinsic atmospheric $\\nu_{\\tau}$ flux, we extend the validity of a previous calculation from $E\\leq 10^{6}$ GeV up to $E \\leq 10^{11}$ GeV. The oscillated atmospheric $\\nu_{\\tau}$ flux is, on the other hand, rather suppressed. We find that, for $10^{3} \\leq E/\\mbox{GeV} \\leq 5\\cdot 10^{7}$, the oscillated $\\nu_{\\tau}$ flux along the galactic plane dominates over the maximal intrinsic atmospheric $\\nu_{\\tau}$ flux, i.e., the flux along the horizontal direction. We also briefly mention the presently envisaged prospects for observing these high-energy tau neutrinos. ", + "introduction": "Searching for high-energy tau neutrinos ($E \\geq 10^{3}$ GeV) will yield quite useful information about the highest energy phenomenon occurring in the universe \\cite{Halzen:2001ty}. The same search may also provide evidence for physics beyond the standard model \\cite{Fukuda:1998mi}. The latter is suggested by the recent measurements of the atmospheric muon neutrino deficit, though yet there is no observation of oscillated atmospheric tau neutrinos at a significant confidence level \\cite{Fukuda:2000np}. Interestingly, it is also only recently that we have the first evidence of existence of tau neutrinos \\cite{Kodama:2000mp}. The high-energy tau neutrinos can be produced in $pp$ and $p\\gamma $ interactions taking place in cosmos. These interactions produce unstable hadrons that decay into tau neutrinos. In this paper, we mainly concentrate on $pp$ interactions and will only briefly comment on $p\\gamma $ interactions as source interactions for producing high-energy tau neutrinos. There can be several astrophysical sites where the $pp$ interactions may occur. Examples of these include the relatively nearby and better known astrophysical sites such as our galaxy and the earth atmosphere, where the basic $pp$ interactions occur as $pA$ interactions. The $pp$ interactions in these sites form a rather certain background to the extra-galactic high-energy tau neutrino searches. It is possible that such interactions are the only sources of high-energy tau neutrinos should the search for high-energy tau neutrinos originating from several proposed distant sites such as AGNs, GRBs, as well as groups and clusters of galaxies, turns out to be negative. Therefore, it is rather essential to investigate the high-energy tau neutrino flux expected from our galaxy and the earth atmosphere. Both the galactic and atmospheric tau neutrinos can be categorized into intrinsic and oscillated ones. Here, intrinsic $\\nu_{\\tau}$ flux refers to the $\\nu_{\\tau}$ produced directly by an interaction while oscillated $\\nu_{\\tau}$ refers to the $\\nu_{\\tau}$ resulted from the $\\nu_{\\mu} \\to \\nu_{\\tau}$ oscillation. Presently, there exists no estimate for the intrinsic high-energy $\\nu_{\\tau}$ flux originating from our galaxy in $pp$ interactions, although the estimates for $\\nu_e$ and $\\nu_{\\mu}$ fluxes due to the same interactions are available \\cite{Stecker:1978ah}. In this work, we calculate the intrinsic $\\nu_{\\tau}$ flux from our galaxy by using the perturbative and nonperturbative QCD approaches to model the $pp$ interactions, and taking into account all major tau neutrino production channels up to $E\\leq 10^{11}$ GeV. We note that the production of tau neutrinos in the terrestrial context was discussed in Ref.~\\cite{DeRujula:1992sn}, which uses a nonperturbative QCD approach for $pp$ interactions. To calculate the oscillated galactic and atmospheric $\\nu_{\\tau}$ fluxes, we apply the two-flavor neutrino oscillation analysis \\cite{Fukuda:2000np}. It is essential to compare the above galactic tau neutrino flux with the flux of atmospheric tau neutrinos. The intrinsic atmospheric tau neutrino flux has been calculated for $E\\leq 10^6$ GeV \\cite{Pasquali:1998xf}. In this work, we extend the calculation up to $E\\leq 10^{11}$ GeV. Such an extension requires the input of cosmic-ray flux spectrum for an energy range beyond that considered in Ref.~\\cite{Pasquali:1998xf}. Furthermore, for a greater neutrino energy, the solutions of cascade equations relevant to the neutrino production behave differently. For the oscillated $\\nu_{\\tau}$ flux, it is interesting to note that the oscillation length for $\\nu_{\\mu}\\to \\nu_{\\tau}$ for the energy range $10^{3} \\leq E/\\mbox{GeV} \\leq 10^{11}$ is much greater than the thickness of the earth atmosphere. Hence the oscillated atmospheric $\\nu_{\\tau}$ flux in this case is highly suppressed. One may argue that the interaction of the high-energy cosmic-rays with the ubiquitous cosmic microwave background (CMB) photons present in cosmos ($p\\gamma $ rather than $pp$ interactions) could also be an important source for high-energy astrophysical tau neutrinos. The center-of-mass energy ($\\sqrt{s}$) needed to produce a $\\tau$ lepton and a $\\nu_\\tau $ is at least $\\sim $ 1.8 GeV. In a collision between a proton with an energy $E_p$ and a CMB photon with an energy $E_{\\gamma_{\\rm CMB}}$, the center-of-mass energy squared of the system satisfies $m_p^2 < s < 4 E_p E_{\\gamma_{\\rm CMB}} + m_p^2$. Since the peak of the CMB photon flux spectrum with a temperature $\\sim $ 2.7 K is at about $2.3\\cdot 10^{-4}$ eV, it requires a very energetic proton with $E_p \\agt 2.5\\cdot 10^{12}$ GeV in order to produce a $\\tau \\nu_\\tau$ pair. Thus, the contribution of the intrinsic tau neutrino flux from the interaction between the cosmic proton and the CMB photon is negligible, unless we are considering extremely high-energy protons, beyond the presently observed highest energy cosmic-rays \\cite{Nagano:ve}. To compute the oscillated $\\nu_{\\tau}$ flux in this case, we consider the non-tau neutrino flux generated by the $p\\gamma$ interaction via the $\\Delta $ resonance, commonly referred to as the Greisen-Zatsepin-Kuzmin (GZK) cutoff interaction, which assumes that the proton travels a cosmological distance \\cite{Greisen:1966jv}. A recent calculation of the intrinsic non-tau GZK neutrino flux indicates that this flux peaks typically at $E\\sim 10^{9}$ GeV \\cite{Engel:2001hd}, beyond the reach of presently operating high-energy neutrino telescopes such as AMANDA and Baikal experiments \\cite{Halzen:2001ty}. This flux decreases for $E< 10^{9}$ GeV. In fact, it falls below the intrinsic non-tau galactic-plane neutrino flux for $E\\leq 5\\cdot 10^{7}$ GeV. The neutrino flavor oscillations of the intrinsic non-tau GZK neutrinos into tau neutrinos result in a $\\nu_{\\tau}$ flux comparable to the original non-tau neutrino flux \\cite{Athar:2000je}. Therefore, in the absence (or smallness) of tau neutrino flux from other possible extra-galactic astrophysical sites, the only source of high-energy tau neutrinos besides the atmospheric background is from our (plane of) galaxy, typically for $10^{3} \\leq E/\\mbox{GeV} \\leq 5\\cdot 10^{7}$. This is an energy range to be explored by the above high-energy neutrino telescopes in the near future. The organization of the paper is as follows. In Section II, we discuss the calculation of intrinsic high-energy tau neutrino flux from our galaxy, including the description of the flux formula, the galaxy model, and the various tau neutrino production channels taken into account in the calculation. Although this flux will be shown small, we shall go through some details of the calculation since they are also useful for the calculation in the next section. In Section III, we present our result on the intrinsic atmospheric $\\nu_{\\tau}$ flux and compare it with the galactic one. In Section IV, we discuss the effects of neutrino flavor mixing, which are used to construct the oscillated $\\nu_{\\tau}$ fluxes from the galaxy and the earth atmosphere respectively. The total galactic $\\nu_{\\tau}$ flux (the sum of intrinsic and oscillated fluxes) is compared to its atmospheric counterpart, and the dominant energy range for the former flux is identified. We also mention the currently envisaged prospects for identifying the high-energy tau neutrinos. We summarize in Section V. ", + "conclusions": "We have calculated the $\\nu_{\\tau}$ flux due to $pp$ interactions in our galaxy. This flux consists of intrinsic tau neutrino flux and that arising from the oscillations of muon neutrinos. We note that the latter flux is dominant over the former by four to five orders of magnitude for the considered neutrino energy range. From Fig. 4, one can see that the main background for the search of high-energy extra-galactic tau neutrino is due to the muon neutrinos produced in the galactic-plane, which then oscillate into tau neutrinos. Such a flux dominates for $10^{3} \\leq E/\\mbox{GeV} \\leq 5\\cdot 10^{7}$. Therefore it is clear that searching for extra-galactic tau neutrinos orthogonal to the galactic-plane is more prospective. In the calculation of galactic tau neutrino flux, we have used a simplified model of matter distribution along our galactic plane to obtain the maximal intrinsic tau neutrino flux. We have explicitly calculated the contribution of heavier states such as $b\\bar{b}$, $t\\bar{t}$ as well as $W^{*}$ and $Z^{*}$ in addition to the more conventional $D_{s}$ channel to the intrinsic tau neutrino flux. We have estimated the average fraction of the incident cosmic-ray energy that goes into tau neutrinos and found it to be less than 1$\\%$. The contributions from $b\\bar{b}$, $t\\bar{t}$, $W^{*}$ and $Z^{*}$ channels are comparable to $D_s$ for $E \\geq 10^{9}$ GeV. For $D_{s}$ channel, we have used both perturbative and nonperturbative QCD approaches. We have extended a previous calculation of intrinsic atmospheric $\\nu_{\\tau}$ flux from $E \\leq 10^{6}$ GeV up to $E \\leq 10^{11}$ GeV. Here, we used the nonperturbative QCD approach to calculate the production of $D_{s}$ mesons in $pA$ interactions. In comparison with the intrinsic galactic-plane $\\nu_{\\tau}$ flux, it is large. However, since the distance between the detector and the neutrino source in the galactic plane is sufficiently large, the neutrino flavor oscillations of non-tau neutrinos into tau neutrinos makes the eventual tau neutrino flux along the galactic plane greater than the atmospheric tau neutrino flux for $10^{3} \\leq E/\\mbox{GeV} \\leq 5\\cdot 10^{7}$. However, the intrinsic atmospheric $\\nu_{\\tau}$ flux dominates over the oscillated galactic $\\nu_{\\tau}$ flux in the direction orthogonal to the galactic plane. We have also briefly mentioned the presently envisaged prospects for observations. In summary, we have completed the compilation of all definite sources of tau neutrino flux, i.e., those from our galaxy and from the earth atmosphere. Such a compilation is needed before one conducts the search for tau neutrinos from extra-galactic sources." + }, + "0112/hep-ph0112014_arXiv.txt": { + "abstract": "The Ultra High Energy Cosmic Ray (UHECR), by $\\nu_r$-Z showering in Hot Dark Halos (HDM), shows an energy spectra, an anisotropy following the relic neutrino masses and clustering in dark halo. The lighter are the relic $\\nu$ masses, the higher their corresponding Z resonance energy peaks. A {\\em twin }light neutrino mass splitting may reflect a {\\em twin } Z resonance and a complex UHECR spectra modulation ({\\em a twin }) bump at highest GZK energy cut-off. Each possible $\\nu$ mass associates a characteristic dark halo size (galactic, local, super cluster) and its anisotropy due to our peculiar position within that dark matter distribution. The expected Z or WW,ZZ showering into $p$ $\\bar{p}$ and $n$ $\\bar{n}$ should correspond to peculiar clustering in observed UHECR at $10^{19}$,$2\\cdot 10^{19}$, $4 \\cdot 10^{19}$. A $\\nu$ HDM halo around a Mpc will allow to the UHECR $n$ $\\bar{n}$ secondary component at $E_n> 10^{20}$ eV (due to Z decay) to arise playing a role comparable with the charged $p$ $\\bar{p}$ ones. Their un-deflected $n$ $\\bar{n}$ flight is shorter leading to a prompt and hard UHECR trace pointing toward the original UHECR source direction. The direct $p$ $\\bar{p}$ pairs are split and spread by random magnetic fields into a more diluted and smeared and lower energy UHECR signal around the original source direction. Additional prompt TeVs signals by synchrotron radiation of electro-magnetic Z showering must also occur solving the Infrared-TeV cut-off. The observed hard doublet and triplets spectra, their time and space clustering already favor the rising key role of UHECR $n$ $\\bar{n}$ secondaries originated by $\\nu$-Z tail shower. ", + "introduction": "Light Neutrino may play a relevant role in Hot Dark Matter models within a hot-cold dark matter (HCDM) scenario. Their clustering in Galactic, Local dark halos offer the possibility to overcome the Cosmic Black Body opacity ($\\gtrsim 4 \\cdot 10^{19}\\,eV$) (GZK) at highest energy cosmic ray astrophysics. These rare events almost in isotropic spread are probably originated by blazars AGN, QSRs in standard scenario, and they should not come, if originally of hadronic nature, from large distances (above tens Mpcs) because of the electro-magnetic dragging friction of cosmic 2.75 K BBR and of the inter-galactic radio backgrounds (GZK cut-off). Indeed as noted by Greisen, Zatsepin and Kuzmin (K.Greisen 1966, Zat'sepin 1966), proton and nucleon mean free path at E $> 5 \\cdot 10^{19} \\,EeV$ is less than 30 $Mpc$ and asymptotically nearly ten $Mpc$.; also gamma rays at those energies have even shorter interaction length ($10 \\,Mpc$) due to severe opacity by electron pair production via microwave and radio background interactions (R.J.Protheroe 1997). Nevertheless these powerful sources (AGN, Quasars, GRBs) suspected to be the unique source able to eject such UHECRs, are rare at nearby distances ($\\lesssim 10 \\div 20 \\, Mpc$, as for nearby $M87$ in Virgo cluster); moreover there are not nearby $AGN$ in the observed UHECR arrival directions. Strong and coherent galactic(R.J.Protheroe 1997) or extragalactic (Farrar et all. 2000) magnetic fields, able to bend such UHECR (proton, nuclei) directions are not really at hand. The needed coherent lengths and strength are not easily compatible with known cosmic magnetic fields. Finally in this scenario the $ZeV$ neutrons born, by photo-pion proton conversions on BBR, may escape the magnetic fields bending and should keep memory of the arrival direction, leading to (unobserved) clustering toward the primary source. Secondaries EeV photons (by neutral pion decays) should also abundantly point and cluster toward the same nearby $AGN$ sources (P.Bhattacharjee 2000, Elbert et all. 1995) contrary to $AGASA$ data. Another solution of the present GZK puzzle, the Topological defects ($TD$), assumes as a source, relic heavy particles of early Universe; they are imagined diffused as a Cold Dark Matter component, in galactic or Local Group Halos. Nevertheless the $TD$ fine tuned masses and ad-hoc decays are unable to explain the growing evidences of doublets and triplets clustering in $AGASA$~ $UHECR$ arrival data. On the other side there are growing evidences of self-correlation between UHECR arrival directions with far Compact Blazars at cosmic distance well above GZK cut-off (Tinyakov P.G.et Tkachev 2001). Therefore the solution of UHECR puzzle based on primary Extreme High Energy (EHE) neutrino beams (from AGN) at ZeV $E_{\\nu} > 10^{21}$ eV and their undisturbed propagation from cosmic distances up to nearby calorimeter (made by relic light $\\nu$ in dark galactic or local dark halo (Fargion et Salis 1997, Fargion, Mele et Salis 1999, Weiler 1999, Yoshida et all. 1998) is still, the most favorite convincing solution for the GZK puzzle. New complex scenarios for each neutrino mass spectra are then opening and important signature of UHECR Z,WW showering must manifest in observed anisotropy and space-time clustering.\\\\ ", + "conclusions": "UHECR above GZK may be naturally born by UHE $\\nu$ scattering on relic ones. The target cosmic $\\nu$ may be light and dense as the needed ones in HDM model (few eVs). Then their $W^+ W^-,ZZ$ pair productions channel (not just the Z resonant peak) would solve the GZK puzzle. At a much lighter, but fine tuned case $m_{\\nu}\\sim 0.4 eV$, $m_{\\nu}\\sim 1.5 eV$ assuming $E_{\\nu}\\sim 10^{22} eV$, one is able to solve at once the known UHECR data at GZK edge by the dominant Z peak; in this peculiar scenario one may foresee (fig.2-3) a rapid decrease (an order of magnitude in energy fluence) above $3\\cdot10^{20}eV$ in future data and a further recover (due to WW,ZZ channels) at higher energies. The characteristic UHECR fluxes will reflect the averaged neutrino-neutrino interactions shown in Fig.2-7. Their imprint could confirm the neutrino masses value and relic density. At a more extreme lighter neutrino mass, occurring for $m_{\\nu}\\sim m_{\\nu_SK}\\sim 0.05 eV$, the minimal $m_{\\nu_{\\tau}},m_{\\nu_{\\mu}}$ small mass differences might be reflected, in a spectacular way, into UHECR modulation quite above the GZK edges. The \"twin\" lightest masses (Fig.5-6-7) call for either gravitational $\\nu$ clustering above the expected one or the presence of relativistic diffused background. Possible neutrino gray body spectra, out of thermal equilibrium, at higher energies may also arise from non standard early Universe. The UHECR acceleration is not yet solved, but their propagation from far cosmic volumes is finally allowed. The role of UHE neutrons in Z-showering, their directional flight leading to clustering in self collimated data is possibly emerging by harder spectra. Peculiar secondaries of TeVs tails may be precursor and afterglows signal correlated to past or future UHECRs pointing toward the same far sources. The IR-TeV solution may be just be a necessary corollary of the Z-Showering GZK solution (Fargion, Grossi et Lucentini 2001). The time and space directional may be a new fundamental test of present Z-Showering model. The discover of UHE neutrino at GZK energies might be testify on ground by UHE $\\tau$ air-shower, born by direct $10^{19}$eV UHE $\\nu$ crossing small Earth crust depth, flashing from the horizontal edges to mountain,balloon and satellite detectors (Fargion 2000). The new generation UHECR data within next decade, may also offer the probe of lightest elementary particle masses, their relic densities, their spatial map distribution and energies and the most ancient and evasive shadows of earliest $\\nu$ cosmic relic backgrounds." + }, + "0112/astro-ph0112301_arXiv.txt": { + "abstract": "A limited-area survey of the $\\eta$ Chamaeleontis cluster has identified 2 new late-type members. The more significant of these is ECHA J0843.3--7905 (= {\\it IRAS\\,} F08450--7854), a slowly-rotating ($P = 12$ d) M2 classical T Tauri (CTT) star with a spectrum dominated by Balmer emission. At a distance of 97 pc and cluster age of $\\approx 9$ Myr, the star is a nearby rare example of an `old' CTT star and promises to be a rewarding laboratory for the study of disk structure and evolution in pre-main sequence (PMS) stars. The other new member is the M4 weak-lined T Tauri (WTT) star ECHA J0841.5--7853, which is the lowest mass ($M \\approx 0.2$ M$_{\\odot}$) primary known in the cluster. ", + "introduction": "\\subsection{Spectroscopy of the new cluster members} \\subsubsection{ECHA J0841.5--7853} Spectroscopy of ECHA J0841.5--7853 (Fig. 3) shows the star to be of spectral type $\\approx$ M4. The level of optical activity is typical for a late-M WTT star, with the H$\\alpha$ equivalent width ({\\it EW\\,}) = $-12$ \\AA. $\\lambda 6707$ Li I absorption is strong ({\\it EW\\,} = 0.9 \\AA) and the spectrum also shows weak $\\lambda 6300$ [OI] emission. The blue spectrum of the star (not shown) shows weak H$\\beta$ and H$\\gamma$ emission; however the poor S/N ratio of the spectrum prevented reliable measurement of these lines. At $V = 17.1$ and an inferred mass $M \\approx 0.2$ M$_{\\odot}$ from Siess et al. (2000) tracks, the star is the optically faintest and lowest mass primary known in the cluster. \\subsubsection{ECHA J0843.3--7905} The more significant of the 2 new members is ECHA J0843.3--7905. The star is clearly a CTT star with an optical spectrum dominated by strong Balmer and Ca II emission (Fig. 4). The H$\\alpha\\/EW = -110$ \\AA, and the H$\\beta\\/EW = -26$ \\AA. Also present in emission in the blue spectrum (Fig. 4a) is $\\lambda 4063$ Fe I, $\\lambda 4471$ He I and all three of the chromospheric lines of the Fe II (42) multiplet ($\\lambda 4923$, 5018 and 5169). The red spectrum (Fig. 4b) shows strong $\\lambda 6300$, 6363 [OI] emission, and emission lines of [N II], He I (the $\\lambda 6678$ transition is shown; also the $\\lambda 7065$ transition is present) and [S II]. $\\lambda 6707$ Li I is present in absorption with {\\it EW\\,} = 0.6 \\AA. Consideration of the spectrum of the star yields a spectral type of M2$-$M3. The colour of the star ($V-I \\approx 2.2$) yields a similar result ($\\approx$ M2). This classification is tentative given the likely high level of optical veiling in the spectrum. The star also has a high degree of optical variability, which we discuss below. From the colour-magnitude diagram and Siess et al. (2000) models, the inferred mass of the star is $M = 0.3-0.4$ M$_{\\odot}$. \\begin{figure} \\begin{center} \\epsfxsize=8.4cm \\epsffile{cttf3.eps} \\caption{Spectrum of the WTT star ECHA J0841.5--7853 near H$\\alpha$, shown here at reduced resolution. Key emission and absorption lines in the spectrum are identified. The spectrum is also shown scaled by a factor of 10 to highlight weak emission and absorption features.} \\end{center} \\end{figure} \\begin{figure} \\begin{center} \\epsfxsize=8.4cm \\epsffile{cttf4.eps} \\caption{(a) Blue and (b) red spectrum of the CTT star ECHA J0843.3--7905, shown here at reduced resolution. Key emission and absorption lines are identified. In (b), the spectrum is also shown scaled by a factor of 10 to highlight weak emission features and the presence of $\\lambda 6707$ Li I.} \\end{center} \\end{figure} \\subsection{Multi-epoch {\\it V\\,} observations} Multi-epoch observations were made of ECHA J0841.5--7853 and ECHA J0843.3--7905 during 1999 February--March and 2000 February--March as part of a larger study to measure the rotation periods of the late-type RECX stars in the $\\eta$ Cha cluster. Lawson et al. (2001) found that all of these stars were variable, with periods attributed to rotational modulation of cool starspots. Differential $V$-band observations of the 2 new cluster members were obtained with respect to nearby stars within their CCD frames. The 1999 and 2000 data sets for each star were independently analysed using the Lomb-Scargle Fourier method for non-equally spaced data, with the frequency range of $f = 0-2$ d$^{-1}$ examined for periodicities. Both stars showed periodic variations in each year. Phased light curves are shown in Fig. 5, and details of the periodicities are listed in Table 2. The S/N ratio for each periodicity was determined by measuring the residual noise level in the pre-whitened data sets. For ECHA J0841.5--7853, a low-amplitude 1.73-d periodicity was recovered in each year. As with the late-type RECX stars, we associate the variations with the rotational modulation of cool spots. Following Allain et al. (1996) we calculated the lower limit to the fractional spot coverage by assuming the spots were dark. The spot fractions are listed as percentages in Table 2. Lawson et al. (2001) corrected the {\\it VRI\\,} photometry of the RECX stars for the effects of the starspots by estimating the $V$ mag and colours of the unspotted, or minimum-spotted, star from the phase of the {\\it VRI\\,} photometry and the $V$-band amplitude of the star. The {\\it VRI\\,} measurements listed in Table 1 were obtained near maximum light and thus represent the unspotted, or minimum-spotted, star. For ECHA J0843.3--7905, a 12-d periodicity was measured in each year, but with different structure in the light curve. During 1999 February, a short-duration ($2-3$ d) peak was observed that may have been an optical flare. Underlying the peak of $\\Delta V = 0.65$ mag is a $0.2-0.3$ mag amplitude quasi-sinusoidal variation that dominates the Fourier analysis. The light curve during 2000 has a different appearance, but with the same underlying period and $V$-band amplitude. The {\\it VRI\\,} photometry reported in Table 1 was obtained in 2000 near maximum light. Given the large photometric amplitude and the variable structure of the light curve probably driven by accretion hotspots, these data may not represent the unspotted photospheric values. \\begin{table} \\centering \\caption{Periods present in the $V$-band light curves during 1999 and 2000. The final column gives lower limits to the dark spot coverage for ECHA J0841.5-7853. ECHA J0843.3-7905 is assumed to have variability due to bright features; see Section 3.2 for details.} \\begin{tabular}{@{}cccccc@{}} \\hline ECHA & Year & Period & Amp. & S/N & Coverage \\\\ & & (d) & (mag) & ratio & (per cent) \\\\ \\hline J0841.5--7853 & 1999 & 1.73 & 0.04 & 3 & 4 \\\\ & 2000 & 1.73 & 0.07 & 3 & 6 \\\\ J0843.3--7905 & 1999 & 12.8 & 0.25 & 4 & -- \\\\ & 2000 & 12.2 & 0.44 & 5 & -- \\\\ \\hline \\end{tabular} \\end{table} \\begin{figure} \\begin{center} \\epsfxsize=8.4cm \\epsffile{cttf5.eps} \\caption{Phase-folded $V$ light curves for ECHA J0841.5--7853 and ECHA J0843.3--7905 obtained during 1999 (left panels) and 2000 (right panels). Within each panel, the period and peak-to-peak amplitude determined from the Fourier analysis of the light curves is shown. Ordinate tick marks are separated by 0.1 mag in all cases. These results are discussed in Section 3.2.} \\end{center} \\end{figure} \\subsection{Comments on the rejected candidates} Analysis of the MSSSO spectra showed 3 of these stars (GSC 9402\\_1003, GSC 9398\\_0099 and GSC 9403\\_0831) were late-K or early-M spectral type giants with H$\\alpha$ in narrow absorption with {\\it EW\\,} $\\approx 1$ \\AA\\, and no detectable $\\lambda 6707$ Li I absorption (see Table 1). All 3 stars can also be ruled out as cluster members from their proper motions. GSC 9398\\_0099 and GSC 9403\\_0831 have Tycho-2 (H$\\o$g et al. 2000) proper motions ($\\mu_{\\alpha}$, $\\mu_{\\delta}$) = (3.7, 1.4) and ($-9.5$, 6.2) mas\\,yr$^{-1}$, respectively (1$\\sigma$ uncertainties are $3-4$ mas\\,yr$^{-1}$). GSC 9402\\_1003 has a USNO UCAC1 (Zacharias et al. 2000) proper motion of ($\\mu_{\\alpha}$, $\\mu_{\\delta}$) = ($-0.3$, 25.3) mas\\,yr$^{-1}$ (1$\\sigma$, 12.5 mas\\,yr$^{-1}$). These values differ by $> 2 \\sigma$ from the {\\it Hipparcos\\,}/Tycho-2 proper motion for the cluster of ($\\mu_{\\alpha}$, $\\mu_{\\delta}$) = ($-30.0 \\pm 0.3$, $27.8 \\pm 0.3$) mas\\,yr$^{-1}$ (Mamajek et al. 2000)\\footnote{Note that Table 2 of Mamajek et al. (2000) incorrectly lists $\\mu_{\\delta} = -27.8$ mas\\,yr$^{-1}$.}. The remaining candidate is listed as `Anonymous' in Table 1, since we found no catalogue entries for this star. Both the spectrum and the colours of the star are consistent with a M5 spectral classification. The star has weak H$\\alpha$ emission and no detectable $\\lambda 6707$ Li I absorption line. Study of on-line scanned plates used to compile the USNO-A2.0 catalogue (Monet et al. 1998) show the star has high proper motion (see Fig. 6). We determined the position of the star against USNO-A2.0 positions for several nearby stars at 4 epochs (1978.10, 1986.18, 1996.13 and 2000.12); the first 3 from scanned plates available from USNO, and the last from analysis of a SAAO CCD image. From these positions we derived ($\\mu_{\\alpha}$, $\\mu_{\\delta}$) $\\approx$ ($-140$, 360) mas\\,yr$^{-1}$. The USNO-A2.0 catalogue rejects stars with $\\mu > 300$ mas\\,yr$^{-1}$. The star is likely a dMe. If it is a main-sequence star, then the star has a distance of $\\sim 50$ pc and tangential velocity of $\\sim 100$ km\\,s$^{-1}$. \\begin{figure} \\begin{center} \\epsfxsize=8.0cm \\epsffile{cttf6.eps} \\caption{USNO finder charts (width = 2.5 arcmin) for the `Anonymous' candidate, obtained from Schmidt plates taken at epochs (a) 1978.10, (b) 1986.18 and (c) 1996.13. The plate source (SRC = Science and Engineering Research Council, ESO = European Southern Observatory, AAO = Anglo-Australian Observatory), emulsion ($J$ or $R$) and plate number is given.} \\end{center} \\end{figure} ", + "conclusions": "Hillenbrand \\& Meyer (1999) examined the frequency of disks as a function of stellar age for nearby clusters and star forming regions and found that disks disperse on a timescale of $< 10$ Myr, with few disks remaining at ages $> 15$ Myr. Haisch, Lada \\& Lada (2001) conducted a similar study on the stellar population of several young clusters, concluding that essentially all stars lose their disks within $\\sim 6$ Myr. The prevalence of disks in the $\\approx 9$ Myr-old $\\eta$ Cha cluster will be examined by Lyo et al. (in preparation). An `old' and nearby CTT star such as ECHA J0843.3--7905 is therefore rare. A similar star is the CTT star TW Hya, also nearby ($d = 56$ pc) and $\\approx 10$ Myr old (Webb et al. 1999). The resolved pole-on disk surrounding TW Hya has been a focus for study of disk structure (e.g. Trilling et al. 2001) and the early planet formation environment. ECHA J0843.3--7905 is of later spectral type and lower mass than TW Hya. Imaging studies of ECHA J0843.3--7905 could give valuable insight into the nature of evolved disks (and indirectly planets) around dwarf M stars. Like other CTT stars, ECHA J0843.3--7905 has a strong infrared excess. Analysis of $L$-band imaging of the cluster core obtained with the South Pole Infrared Explorer telescope during 1999 (Lyo et al., in preparation) found $L = 7.8$ for this star. Assuming an M2 spectral type, the $L$-band excess is $\\approx 2$ mag. The {\\it IRAS\\,} FSC entry for the star indicates high-quality 25- and 60-$\\mu$m fluxes. Our survey of $\\approx 40$\\% of the known extent of the cluster found 2 new cluster members not detected by the discovery {\\it ROSAT\\,} HRI image of Mamajek et al. (1999, 2000), thereby increasing the number of stellar primaries to 15. A survey of similar depth across the cluster might therefore find only several more new members. This result appears to be at odds with Mamajek et al. (2000) who predicted from consideration of the cluster IMF that the stellar population was $2-4 \\times$ the (then) known number of 13 primaries. Our survey might indicate that the cluster extent is not constrained by the {\\it ROSAT\\,} HRI field and that $15-40$ primaries await discovery beyond the HRI boundary. Alternatively, if the cluster is constrained by the HRI field, then the low success rate of our study suggests the cluster may contain as few as $\\approx 20$ primaries. (These estimates do not address the brown dwarf population expected to accompany the stellar members.) Either of the above population scenarios indicates the X-ray survey must have been relatively complete at detecting cluster members within the HRI field. This result confirms the unusual skewness of the ratio of X-ray luminosity to bolometric luminosity noted by Mamajek et al. (1999, 2000), with most of the late-type RECX stars having a flux ratio near the `saturation' level of log $L_{X}$/$L_{\\rm bol} \\approx -3$." + }, + "0112/astro-ph0112137_arXiv.txt": { + "abstract": "We have studied the 1999 soft X-ray transient outburst of XTE J1859+226 at radio and X-ray wavelengths. The event was characterised by strong variability in the disc, corona and jet -- in particular, a number of radio flares (ejections) took place and seemed well-correlated with hard X-ray events. Apparently unusual for the {\\em canonical} `soft' X-ray transient, there was an initial period of low/hard state behaviour during the rise from quiescence but {\\em prior} to the peak of the main outburst -- we show that not only could this initial low/hard state be an ubiquitous feature of soft X-ray transient outbursts but that it could also be extremely important in our study of outburst mechanisms. ", + "introduction": "\\begin{figure*} \\begin{center} \\leavevmode \\psfig{file=lightcurve-xrad.ps,height=22cm,angle=0} \\caption{Radio lightcurve from the VLA, MERLIN, the Ryle Telescope, the RATAN Telescope and a few additional points from the GBI plotted with the hard and soft X-ray data from {\\sl CGRO}/BATSE and {\\sl RXTE}/ASM.} \\label{lightcurve} \\end{center} \\end{figure*} \\begin{figure} \\begin{center} \\leavevmode \\psfig{file=merlin.eps,height=8cm, angle=0} \\caption{The 1.66 GHz lightcurve of MERLIN, showing the excellent coverge of the first and major flare.} \\label{merlin} \\end{center} \\end{figure} X-ray transients are a group of X-ray binaries which exhibit dramatic outbursts, generally at all wavelengths, due to changes in mass accretion rate. A `typical' soft X-ray transient event (or `X-ray nova') will show a fast rise to its maximum luminosity, accompanied by X-ray spectral softening, and reach the high or sometimes the very high state; this is then followed by a slow exponential decay back to quiescence (see M\\'endez et al. 1998). More recently fewer transients have appeared to behave in this typical manner, almost certainly due (in part?) to our improved observing coverage. Some show a period of hard state behaviour before reaching the full outburst (e.g. XTE J1550$-$564 Wilson \\& Done 2001); others do not progress beyond the low/hard state at all (the hard state transients -- see e.g. Brocksopp et al. 2001 and references within). Note that the terms `low' and `high' states do not necessarily refer to the flux of a source but instead indicate the relative contributions of hard and soft X-ray flux -- see van der Klis (1995) and Nowak (1995) for full details of spectral state classification. XTE J1859+226 was discovered by the All Sky Monitor onboard the Rossi X-ray Timing Explorer satellite on 1999 October 9 (MJD 51460; Wood et al. 1999). On its initial detection the soft (2--12 keV) X-ray flux was $\\sim$ 160 mCrab and rising at $\\sim$ 6 mCrab/hour. Follow-up PCA observations revealed a hard power law energy spectrum and a power spectrum with a 0.45 Hz quasi-periodic oscillation (QPO), plus harmonics (Markwardt et al. 1999). Hard X-ray observations by BATSE (of the Compton Gamma Ray Observatory) confirmed the hard spectrum up to 200 keV and showed that the hard X-ray flux peaked while the soft X-ray source was still rising (McCollough \\& Wilson 1999). We note that this `initial low/hard state' is not thought to be a common property of a `canonical' soft X-ray transient outburst, although we show in Section 4 that it has indeed been observed many time previously. These X-ray properties suggested a likely black hole X-ray binary nature for the system. Despite the intially hard spectrum, suggesting that XTE J1859+226 may be an addition to the list of hard state X-ray transients (Brocksopp et al. 2001; Brocksopp, Bandyopadhyay, Fender 2001, in prep.), as the hard X-ray flux decreased the soft X-ray source entered a series of flares; the source softened and reached a peak of $\\sim$ 1.5 Crab (2--12 keV) about eight days after the initial detection (Focke et al. 2000). PCA observations durin g this flaring period again showed the presence of QPOs, this time at 6--7 Hz and a second at 82--187 Hz (Cui et al. 2000). The radio counterpart was discovered at a flux density of $\\sim$ 10 mJy on 1999 October 11 (MJD 51462) at the Ryle Telescope and VLA (Pooley \\& Hjellming 1999). The optical counterpart was also discovered a day later with an $R$-band magnitude of $\\sim$ 15.1 (Garnavich et al. 1999). Spectroscopy of this source revealed a strong blue continuum with weak emission lines at H$\\alpha$, H$\\beta$, He\\,{\\sc ii}\\,$\\lambda 4686$ and C\\,{\\sc iii}/N\\,{\\sc iii}\\,$\\lambda 4640--4650$ (Garnavich et al. 1999, Wagner et al. 1999). Studies of an interstellar absorption feature suggested a low value ($E(B-V)\\sim 0.58$) for the interstellar extinction (Wagner et al. 1999, Hynes et al. 1999). Following decay of the outburst, attempts have been made to search the optical photometry for the orbital period of the system and a potential modulation at 9.15 hours discovered (Garnavich \\& Quinn 2000); this was confirmed by Sanchez-Fernandez et al. (2000). More recently a mass function of 7.4 M$_{\\odot}$ has been determined, thus confirming this suggestion that the compact object is a black hole (Filipenko \\& Chornock 2001). A QPO of 0.76 mHz was also detected in optical time series data (Chaty et al. 2001). X-ray and optical observations have classified XTE J1859+226 as a `typical X-ray nova' with a fast rise and exponential decay (Chaty et al. 2001). Comparisons with various other X-ray transients have been drawn -- it entered an optical `mini-outburst' 235 days after the initial maximum, similar to GRO J0422+32 (a {\\em hard} state transient) and its optical behaviour has also been likened to that of A0620$-$00 and GRS 1124$-$68 (Charles et al. 2000). Alternatively, the X-ray behaviour is reminiscent of XTE J1550$-$564 (Cui et al. 2000). In this paper we study the radio and X-ray observations of XTE J1859$-$226. The observations are outlined in Section 2 and the results presented in Section 3. Finally we compare certain features of XTE J1859+226 with those of other sources in Section 4 and discuss the implications of these results in Section 5, before drawing our conclusions in Section 6. \\begin{figure} \\begin{center} \\leavevmode \\psfig{file=map.ps,height=9cm} \\caption{8.64 GHz map of XTE J1859+226 taken by the VLA on 1999 December 1. It appears to be resolved, thus confirming our predictions throughout this paper that XTE J1859+226 is a jet source. Further analysis of the radio maps will be presented in a future paper.} \\label{map} \\end{center} \\end{figure} ", + "conclusions": "We have observed XTE J1859+226 at radio and X-ray wavelengths during its 1999 outburst. The results show that this soft X-ray transient spent a few days in the low/hard state before increasing to its maximum and we suggest that this may be a typical feature of soft X-ray transient outbursts, thus providing a means by which we can study the outburst mechanisms. As the outburst continued we found that a series of radio ejections took place, simultaneously with spectral hardening and apparently correlated with the hard X-ray lightcurve, thus emphasizing the strength of the jet/corona connection. Again, it appears that a series of radio ejections correlated with the corona emission may be typical features of these systems." + }, + "0112/astro-ph0112247_arXiv.txt": { + "abstract": "We discuss possible distortions of the ionization history of the Universe in the model with small scale baryonic clouds. The corresponding scales of the clouds are much smaller than the typical galactic mass scales. These clouds are considered in a framework of the cosmological model with the isocurvature and adiabatic perturbations. In this model the baryonic clouds do not influence on the cosmic microwave background anisotropy formation directly as an additional sources of perturbations, but due to change of the kinetics of the hydrogen recombination . We also study the corresponding distortions of the anisotropy and polarization power spectra in connection with the launched MAP and future PLANCK missions. ", + "introduction": "One of the most important problems of the modern cosmology is the determination of the density and spatial distribution of the baryonic fraction of the matter. There are several sources of information about $\\Omega_b h^2 = \\rho_{b}/\\rho_{cr}$ parameter, where $\\rho_{b}$ and $\\rho_{cr}$ are the present values of the baryonic and critical densities and $h$ is the Hubble constant normalized to 100 $\\kms\\Mpc^{-1}$. Firstly, the baryonic fraction of the matter manifests itself in the well known mass -luminosity relation for galaxies and cluster of galaxies which leads to the following value of the $\\Omega_b h^2$ parameter: $\\Omega_b h^2\\simeq 0.028^{+0.009}_{-0.008} $ (see for the review by Freedman et al. 2001). Another one comes from the confrontation of the Standard Big Bang Nucleosynthesis (SBBN) theory and observational data (see for the review by Fukugita, Hogan \\& Peebles 1998). The corresponding value of the baryonic density from this method is $\\Omega_b h^2=0.019 \\pm 0.001.$ An additional empirical relation between baryonic and dark matter fractions $F_{b,m}=\\Omega_b/\\Omega_m\\simeq 0.1$ at $h=0.65$ comes from X-ray data on clusters of galaxies (Carlberg et al. 1996; Ettori \\& Fabian 1999). For the most popular $\\Lambda$CDM cosmological model with $\\Omega_m\\simeq 0.3$ and $\\Omega_{\\Lambda} \\simeq 0.7$, the corresponding value of the $\\Omega_b h^2$ parameter is $\\sim 0.02$ in agreement with the SBBN predictions. An independent important information about the baryonic fraction of the matter in the Universe comes from the recent CMB experiments such as BOOMERANG (de Bernardis et al. 2000) and MAXIMA-1 (Hanany et al. 2000). Fitting the CMB anisotropy power spectrum to the above mentioned observational data (Tegmark \\& Zaldarriaga 2000; White et al. 2000 and Lesgourgues \\& Peloso 2000) indicates that a baryon fraction parameter should be significantly larger than the SBBN expected value, namely, $\\Omega_b h^2\\simeq 0.03$. However, recently Bond \\& Critteden (2001) show that new BOOMERANG, MAXIMA-1 and DASI data do not contradict to $\\Omega_b h^2=0.022 \\pm 0.004$. It is worth noting that the above mentioned methods of the baryonic fraction density estimation from the CMB and SBBN predictions are based on the simple idea that the distribution of matter (including dark matter particles and baryons) is practically homogeneous for all scales, except some fluctuations leading to the galaxy and large-scale structure formation. Typically, they are assumed to be adiabatic one. One can ask, how sensitive are the CMB data themselves to the presence of the small-- scale baryonic (non-linear) clouds before cosmological recombination and how can they transform the standard schemes of the cosmological parameter extraction from the CMB data? Definitely, this possibility is related to the isocurvature perturbations of the composite fluid which contains baryons, CDM particles, photons and neutrinos at very high redshift $z\\gg 10^3$. There are a lot of modes of perturbations in the composite fluid, which is discussed by Riazuelo \\& Langlos (2000), Bartolo, Matarrese \\& Rioto (2001), Polarski \\& Starobinsky (1994), Abramo \\& Finelli (2001), Bucher, Moodley \\& Turok (2000) and others. The general idea about classification of modes of perturbations is based on a very simple definition of the isocurvature modes. They do not perturb the gravitational potential. This means that the fluctuations of the total matter density $\\rho_{tot}$ are zero (see Burns 2001), \\begin{equation} \\delta\\rho_{tot}=\\sum\\limits_{i=0}^{N}\\rho_{i}\\delta_i + 4\\rho_{\\gamma}(1+R_{\\nu \\gamma}) \\frac{\\delta T}{T}=0, \\label{eq:eq1} \\end{equation} where $\\rho_{i}$ denotes the density of each massive species including baryons and different kinds of the CDM particles, $\\delta_i = \\delta\\rho_{i}/\\rho_{i}$ is the density contrast for each massive component, $R_{\\nu \\gamma}$ is the density ratio between neutrinos $\\rho_{\\nu}$ and black body radiation $\\rho_{\\gamma}$, and $\\delta T/T$ is the CMB temperature perturbations. We would like to point out that in the definition of the isocurvature perturbations in Eq.~(\\ref{eq:eq1}) one can find some peculiar mode (or modes) which compensates the baryonic perturbations potential, i.e., it corresponds to the condition $\\rho_b\\delta_b=- \\rho_x\\delta_x$ for some $x$ component of the CDM particles mixture. We call below this mode as a compensate isocurvature mode (CIM) for the $x$-component of the dark matter particles. If several components of the dark matter particles take part in the CIM formation, we will continue to call them as $x$- component.\\footnote{ As one can see this mode corresponds to $\\delta T/T=0$. This means that the CIM are equivalent to an isotemperature perturbations. Note, that exactly the same mode was described by Abramo \\& Finelli (2001), but for compensation between quintessence scalar field perturbations and some kind of the CDM-particles.} In principle, for the CIM perturbations it is possible to assume that the amplitudes $\\delta_x$ and $\\delta_b$ are less than unity or $|\\delta_x|\\sim 1$ and $|\\delta_b|\\simeq - (\\rho_x/\\rho_b)|\\delta_x|\\gg 1$. One of the most interesting cases corresponds to the model with $\\delta_x\\simeq -1 $, which means that some patches of the cosmological matter do not contain the CDM $x$-particles, but at the same patches there are non-linear clouds of the baryonic matter which compensate the perturbations of the gravitational potential. Below we will assume that a typical mass scale of the CIM perturbations is smaller than the typical galactic mass scale. This means that CIM perturbations do not influence on the CMB anisotropy formation as the additional sources of perturbations, but they can transform the kinetics of the hydrogen recombination. This leads to the transformation of the corresponding $C_l$ power spectrum of the CMB for the adiabatic fluctuations at the scales above a few Mpc. It is necessary to note that the idea about non-homogeneous distribution of the baryonic matter at small scales is not new. The importance of the entropic perturbations in the history of the cosmological expansion was {\\it ad hoc} demonstrated by Doroshkevich, Zel'dovich and Novikov (1967) and Peebles (1967,1994) and recently was generalized taking into account multi-species structure of the cosmological plasma by Gnedin \\& Ostriker (1992), Hogan (1993) \\& Loeb (1993), Peebles \\& Juszkiewich (1998). The possible inhomogeneities of the baryon fraction distribution in the epoch of the nucleosynthesis (Inhomogeneous Big Bang Nucleosynthesis -IBBS) was widely discussed in the literature (see for the review by Jedamzik \\& Rehm 2001) in connection with quark-hadron phase transition. But the typical scales of such kind of peculiarities are extremely small compared to the typical mass scale $ M\\sim 10^5 - 10^6 M_{\\odot}$ for the isocurvature perturbations. Other events or processes have been suggested as possible sources of the isocurvature perturbations partly connected with the baryon re-distribution in the space. For example, cosmic strings and corresponding currents and magnetic fields could generate specific features in the baryonic matter (Malaney \\& Butler 1989). Yokoyama \\& Sato (1991), Dolgov \\& Silk (1993), Polarsky \\& Starobinsky (1994), Novikov, Schmalzing \\& Mukhanov (2000) have shown a few different ways for the generation of the isocurvature perturbation in the framework of the inflation theory. In connection with the above--mentioned problem of the baryonic fraction determination from the cosmological nucleosynthesis and the CMB anisotropy data we dedicate our paper to the re-examination of the models with non-linear sub-horizon-scale (at the epoch of the hydrogen recombination) baryonic clouds with $\\delta_b\\gg 1$. Such inhomogeneities do not manifest in the CMB anisotropy( because of the extremely small scales) but manifest themselves by the transformation of the ionization history of the primeval hydrogen-helium plasma at redshift $z\\sim 10^3$. All these factors should be taken into account in the reconstruction of the ionization history of the Universe especially at the period of the cosmological hydrogen recombination . The reason for importance of the possible very small scale entropy perturbations at the epoch of recombination is connected with the very small mean free path of the Ly-$\\alpha$ photons at $z\\sim 10^3$: $l_{L\\alpha} \\sim 3\\times 10^{10}((1+z)/1000)^{-5/2} (\\Omega_b h^2/0.02)^{-1}$ cm which corresponds to the baryon mass $M_{L\\alpha}=4\\pi/3\\rho_b l_{L\\alpha}^3 \\sim 2\\times10^{-23}((1+z)/1000)^{-9/2} (\\Omega_b h^2/0.02)^{-2}M_{\\odot}$ where $M_{\\odot}$ is the Solar mass. High amplitude baryonic clouds with masses $M\\gg M_{L\\alpha}$ could transform the process of recombination at the beginning and dissipate during recombination up to the crucial masses $M_{diss}\\sim 10^5 M_{\\odot}$ (Liu et al. 2001). We will show that if the typical masses of the clouds $M$ are $M>M_{diss}$ the hydrogen and helium recombination inside and outside clouds goes independently. Due to non-linear dependency of the electronic ionization fraction $x_e$ on the baryon density the hydrogen and helium inside the clouds recombine faster than outside them. Thus the dynamics of the mean value of the electronic ionization fraction $x_e$ which plays a crucial role in the CMB anisotropy and polarization formation decreases slower than, for example, in the uniform model with the mean value of the baryonic fraction of the matter. We will show that in the cloudy baryonic plasma the kinetics of the $H-He^4$ recombination is closer to the delayed recombination model by Peebles et al. (2000) with concrete relation between $\\epsilon_{\\alpha}$ and $\\epsilon_{i}$ parameters of their model and amplitudes of perturbations and the filling factor of the baryonic clouds. We will show how sensitive the $C_l$ power spectrum of the CMB anisotropy is to the mentioned above parameters of the baryonic clouds. We will discuss possible manifestation of the small scale perturbations in the MAP and upcoming PLANCK observational data. ", + "conclusions": "As it was mentioned in Introduction, the baryonic fraction of the matter is one of the most important cosmological parameters, which determines the most preferable model of the Universe. Two crucial parameters of the theory are now under discussion, i.e., the light chemical elements abundance, which is related directly to the SBBN predictions, and the $\\Omega_b$ parameter which can be determined from the current and future CMB observations. In both cases the determination of the parameters depends on the hypothesis about the spatial distribution of the baryonic fraction of the matter and can be tested by the CMB experiments such as the launched MAP and the future PLANCK missions. \\begin{figure} \\centering \\epsfig{file=figNN5.eps,width=8.5cm} \\caption{ The functions $D(l)$ for the anisotropy (solid line) and polarization ( dashed line) as functions of the multipole number $l$ for $\\xi=11$ , $f=0.1$ model. } \\label{} \\end{figure} In our paper we have shown how important it is the possible baryonic inhomogeneity at small scales ($M \\ll 10^{13} M_{\\odot}$) in the CMB anisotropy and polarization formation through distortions of the cosmological recombination. The extremal case, when the density contrast $\\xi\\simeq 10$, shows that the clumps in the baryon fraction of the matter can significantly change the amplitudes and positions of the Doppler peaks in the CMB anisotropy and polarization power spectrum. For example in the above mentioned $\\xi\\simeq 10$ model the differences between $C_l$ for the anisotropy and for the polarization in the models with $\\Omega_b h^2=0.02$ ( no clouds) and $\\langle \\Omega_b \\rangle h^2=0.02$ (with clouds) are presented in Fig.5 in a form of the following the functions. For the anisotropy \\begin{equation} D_a(l)=2 (C_{a,nc}(l)-C_{a,c}(l))/(C_{a,nc}(l)+C_{a,c}(l)), \\end{equation} where $C_{a,nc}(l)$ and $C_{a,c}(l)$ denote the non-cloudy and cloudy model, respectively, and the index $a$ corresponds to the anisotropy spectrum. The analogous definition of the $D_p(l)$ function is used in Fig.5 for the CMB polarization. As one can see from Fig.5, practically for all ranges of multipoles the differences between cloudy and non-cloudy models are observable for the PLANCK mission. Moreover, if the parameter $\\xi\\geq 1$, then we need to include the possible cloudy baryonic model in the schemes of the cosmological parameter extractions from the current and future CMB anisotropy and polarization data. As one can see from Eq.(4) and Eq.(5), if $f(\\xi-1)\\ll 1$, then the difference between $\\rho_{b,out}$ and mean baryon density is $\\sim f(\\xi-1) $. According to the predicted accuracy of the cosmological parameter extraction from the PLANCK mission, the corresponding uncertainties for the baryonic fraction of the matter $\\Delta_b= \\delta\\Omega_b/\\Omega_b$ must be less then a few percents. Taking conservative limit $\\Delta_b\\simeq f(\\xi-1)\\sim 0.1$ and $\\xi-1 \\sim 1$ we can obtain that the corresponding fraction $f$ should be detectable by the PLANCK satellite, if $f\\geq 0.1$. We would like to point out that our simple model of the baryonic clouds is based on the one possible modes of the isocurvature perturbations at the small scales. It would be interesting to investigate the more complicated models of the initial perturbations. This program is in progress." + }, + "0112/astro-ph0112071_arXiv.txt": { + "abstract": "We describe a new approach to establishing the mass composition at high energies. Based on measuring both the vertical and inclined shower rates, it has the potential to distinguish heavy nuclei from light nuclei. We apply the method to Haverah Park data above $10^{18}~$eV to show that, under the assumption that the Quark Gluon String Jet Model correctly describes the high energy interactions, the inclined shower measurements favour a light composition at energies above $10^{19}$~eV. The same conclusion is obtained using a variety of assumptions about the cosmic ray spectrum. To the extent that precise spectral measurements will be possible by forthcoming experiments such as the Auger observatories, the method will further constrain data on composition of the ultra high energy cosmic rays. ", + "introduction": "Efforts to understand the origin of cosmic rays at any energy are greatly hampered by our lack of knowledge of the mass distribution in the incoming cosmic ray beam. The determination of the arrival direction of cosmic rays does not depend on knowledge of the mass of the primary and, using the fluorescence technique, the primary energy is obtainable with only a small systematic uncertainty because of the unknown mass. Even with the traditional ground arrays it has been possible to devise ways of deducing the primary energy that are reasonably independent of model and mass uncertainties, at least at the 30 \\% level. However use of the data on the energy spectrum and arrival direction distribution to decide between various origin models does require knowledge of the primary mass distribution. While there is a common assumption that protons dominate at the highest energies, hard experimental evidence is lacking, as it is at energies above 10$^{18}$ eV. Some models of cosmic ray origin lead to the conclusion that photons dominate the cosmic ray beam. This hypothesis can be more readily tested than can competing hypotheses that advocate protons or iron nuclei as the majority primaries. This is because the shower development of a photon primary is very different from that of any hadronic primary. The difference is larger than is expected between the showers created by proton and iron primaries. For example, Halzen {\\it et al.} \\cite{vazquez} have shown that a photon is extremely unlikely to produce a shower in the atmosphere that has a development profile that looks like the famous event of $3 \\times 10^{20}$ eV observed by the Fly's Eye Group \\cite{320EeV}. A study of horizontal showers recorded at Haverah Park has shown that at 10$^{19}$ eV there can be no more than 40 \\% of the primaries that are photons \\cite{PRL}, while a search, using the traditional idea of photon showers being deficient in muons, has given a similar limit at the same energy \\cite{Teshima01}. These analyses all assume that photons $>$ 10$^{19}$ eV have not acquired hadron-like properties, as has sometimes been speculated. There is no reliable evidence on the hadronic mass composition above 10$^{19}$ eV. Some years ago an analysis of the Fly's Eye data \\cite{gaisser} pointed to a change from an iron-dominated composition at $3 \\times 10^{17}$ eV to a proton-dominated composition near 10$^{19}$ eV. This conclusion was drawn from a study of the variation of depth of maximum with energy (the elongation rate) and from an analysis of the spread in depth of maximum at a given energy. The elongation rate behaviour was in agreement with an earlier, model-independent, analysis from Haverah Park \\cite{watsonwalker} and with a related study using Cherenkov light by the Yakutsk group \\cite{Yakutsk}. The conclusions on mass composition from the study of the spread in depth of maximum were confirmed by \\cite{Wibig} but none of these analyses extended to energies above 10$^{19}$ eV. \\begin{figure} \\ybox{0.4}{xmax} % \\caption{Compilation of data on depth of maximum as a function of energy from different experiments compared with predictions for different models. This figure was supplied by D. Heck and J. Knapp and appear in \\cite{heck}} \\label{xmax} \\end{figure} A different conclusion has been reached by the AGASA group based on the variation of the muon content of showers with energy \\cite{MuAgasa}. Their analysis favours a composition that remains 'mixed' over the 10$^{18}$ to 10$^{19}$ eV decade. Furthermore there is a difference between the conclusions reached from the data on depth of maximum from the Fly's Eye experiment \\cite{gaisser} and those reached from the HiRes prototype, operated with the MIA detector, in the range 10$^{17}$ to 10$^{18}$~eV \\cite{HiResMiaDepth}. The more recent data are used to infer protons as early as 10$^{18}$ eV, if the QGSJET model\\cite{QGSJET} is correct. There is a good evidence from Haverah Park data \\cite{compICRC} that the QGSJET model is satisfactory, at least to 10$^{18}$ eV. However the HiRes/MIA conclusion is, in turn, in contradiction with a new analysis of Haverah Park data \\cite{compICRC} that suggests that protons make up only 34 \\% of the flux, independent of the energy, between $3 \\times 10^{17}$ and $2 \\times 10^{18}$ eV, with iron nuclei making up the remainder. The addition of more data in recent years and a better understanding of what might be the optimum shower model has not helped to clarify the position: rather the reverse. In figure \\ref{xmax} we show a compilation of data on depth of maximum, $X_{\\rm m}$, as a function of energy \\cite{heck}. The sensitivity of a conclusion to the choice of model is evident. The matter is thus far from being resolved and there is clearly scope for alternative approaches. ", + "conclusions": "We believe that we have demonstrated the potential of a new method for extracting the mass spectrum above 10$^{18}$ eV by studying cosmic rays at high zenith angles. It is based on the difference in abundance of muons with respect to photons and electrons for different primaries. We are not making strong claims for the finality of our conclusions (see Fig. \\ref{ratio}) about mass composition. A definitive statement awaits a better understanding of the differences between different energy spectra above this energy and an exploration of the sensitivity of the method to different models of particle interactions. This technique can be extended to very high energies and, assuming the interaction model is understood, we would expect to be able to extract the mass spectrum in the range 10$^{19}$ to 10$^{20}$ eV with data from the Pierre Auger Observatory. {\\bf Acknowledgements:} This work was partly supported by a joint grant from the British Council and the Spanish Ministry of Education (HB1997-0175), by Xunta de Galicia (PGIDT00PXI20615PR), by CICYT (AEN99-0589-C02-02) and by PPARC(GR/L40892)." + }, + "0112/hep-ph0112312.txt": { + "abstract": "This is the summary report of the ``Indirect Investigations of SUSY'' subgroup of the P3 Physics Group at Snowmass 2001. ", + "introduction": "\\label{sec:intro} In this report we consider indirect probes of supersymmetry (SUSY). Following our charge, we will review the current experimental status and discuss possible levels of improvements in various measurements sensitive to supersymmetry through either virtual or astrophysics effects. We mention the upcoming experiments which are likely to achieve such precision, and outline which theoretical models and ideas can be tested by those experiments. Since it is impossible to give a detailed review of every single topic here, we have limited our discussion to a few representative topics which were studied either at Snowmass or since then. A large fraction of this report is based on the individual written contributions~\\cite{Feng:2001mq,Feng:2001hm,Ellis:2001us,Feng:2001ut,ge1,ge3} to our subgroup as well as talks presented at Snowmass. In Section~\\ref{sec:gm2} we discuss the implications of the recent $g_\\mu-2$ measurement (and its possible improvements) for supersymmetry. In Section~\\ref{sec:cp} we review searches for \\CP violation and Section~\\ref{sec:LFV} deals with lepton-flavor violation (LFV). Section~\\ref{sec:B} is devoted to the future of B-physics. In Section~\\ref{sec:DMdirect} (\\ref{sec:DMindirect}) we discuss direct (indirect) searches for supersymmetric dark matter. We summarize and present our conclusions in Section~\\ref{sec:discussion}. ", + "conclusions": "\\label{sec:discussion} Fig.~\\ref{fig:reach} shows a compilation of many pre-LHC experiments in astrophysics, as well as particle physics at both the energy frontier and lower energies. The signals considered, the projected sensitivities, and the experiments likely to achieve them, are discussed in \\cite{Feng:2001ut}. On the particle physics side, the signals considered were supersymmetry searches at LEP~\\cite{SUSYWG} and the Tevatron~\\cite{Matchev:1999nb,Baer:1999bq,Barger:1998hp,% Matchev:1999yn,Lykken:1999kp}, the improved measurement of the $B\\rightarrow X_s\\gamma$ branching ratio at B-factories, as well as the projected final sensitivity of the Brookhaven $g_\\mu-2$ experiment. On the astrophysics side, the figure shows the projected reach of the upcoming direct dark matter detection experiments, as well as the multitude of other experiments to detect indirect neutrino, photon or positron signals from neutralino annihilations, as discussed earlier. \\begin{figure}[t] \\includegraphics[height=2.3in]{P3_42_fig13.ps}% \\caption[]{Estimated reaches of various high-energy collider and low-energy precision searches, direct dark matter searches, and indirect dark matter searches before the LHC begins operation, for $\\tb=10$. The projected sensitivities used are given in Ref.~\\cite{Feng:2001ut}. The darker shaded (green) regions are excluded by the requirement that the LSP be neutral (left) and by the LEP chargino mass limit (bottom and right). The regions with potentially interesting values of the LSP relic abundance: $0.025\\le \\Omegachi h^2 \\le 1$ (light-shaded, yellow) and $0.1 \\le \\Omegachi h^2 \\le 0.3$ (medium-shaded, light blue), have also been delineated. The regions probed extend the curves toward the forbidden regions. (From Ref.~\\cite{Feng:2001ut}.)} \\label{fig:reach} \\end{figure} % Several striking features emerge from Fig.~\\ref{fig:reach}. First, we see that, within the minimal supergravity framework, nearly all of the cosmologically preferred models will be probed by at least one experiment\\footnote{While this is strictly true for low $\\tb$, at higher $\\tb$ some of the preferred region may escape all probes, but this requires heavy superpartners and a significant fine-tuning of the electroweak scale~\\cite{Feng:2001zu}. Furthermore, the large $\\tan\\beta$ region is also effectively probed via the \\bs signal~\\cite{Dedes:2001fv} which was not considered here.}. In the most natural regions, all models in which neutralinos form a significant fraction of dark matter will yield some signal before the LHC begins operation. Also noteworthy is the complementarity of traditional particle physics searches and indirect dark matter searches. Collider searches require, of course, light superpartners. High precision probes at low energy also require light superpartners, as the virtual effects of superpartners quickly decouple as they become heavy. Thus, the LEP and Tevatron reaches are confined to the lower left-hand corner, as are, to a lesser extent, the searches for deviations in $B \\to X_s \\gamma$ and $a_{\\mu}$. These bounds, and all others of this type, are easily satisfied in the focus point models with large $m_0$, and indeed this is one of the virtues of these models. However, in the focus point models, {\\em all} of the indirect searches are maximally sensitive, as the dark matter contains a significant Higgsino component. Direct dark matter probes share features with both traditional and indirect searches, and have sensitivity in both regions. It is only by combining all of these experiments, that the preferred region may be completely explored. \\begin{figure}[t] \\includegraphics[height=2.3in]{P3_42_fig14.ps}% \\caption[]{As in Fig.~\\ref{fig:reach}, but in the $(m_0, \\tb)$ plane for fixed $\\mgaugino=400~\\gev$, $A_0 = 0$, and $\\mu > 0$. The regions probed are toward the green regions, except for $\\Phi_{\\gamma}^{50}$, where it is between the two contours. The top excluded region is forbidden by limits on the \\CPn -odd Higgs scalar mass. (From Ref.~\\cite{Feng:2001ut}.)} \\label{fig:reach400} \\end{figure} Finally, these results have implications for future colliders. In the cosmologically preferred regions of parameter space with $0.1 < \\Omegachi h^2 < 0.3$, all models with charginos or sleptons lighter than 300 GeV will produce observable signals in at least one experiment. This is evident for $\\tb=10$ in Fig.~\\ref{fig:reach}. In Fig.~\\ref{fig:reach400}, we vary $\\tb$, fixing $\\mgaugino$ to 400 GeV, which roughly corresponds to 300 GeV charginos. We see that the preferred region is probed for any choice of $\\tb$. (For extremely low $\\tb$ and $m_0$, there appears to be a region that is not probed. However, this is excluded by current Higgs mass limits for $A_0 = 0$. These limits might be evaded if $A_0$ is also tuned to some extreme value, but in this case, top squark searches in Run II of the Tevatron \\cite{Demina:1999ty} will provide an additional constraint.) These results imply that if any superpartners are to be within reach of a 500 GeV lepton collider, some hint of supersymmetry must be seen before the LHC begins collecting data. This conclusion is independent of naturalness considerations. While our quantitative analysis is confined to minimal supergravity, we expect this result to be valid more generally. For moderate values of $\\tb$, if the dark matter is made up of neutralinos, they must either be light, Bino-like, or a gaugino-Higgsino mixture. If they are light, charginos will be discovered. If they are Bino-like, light sfermions are required to mediate their annihilation, and there will be anomalies in low energy precision measurements. And if they are a gaugino-Higgsino mixture, at least one indirect dark matter search will see a signal. For large $\\tb$, low energy probes become much more effective and again there is sensitivity to all superpartner spectra with light superpartners. Thus it appears, on qualitative grounds, that all models in which the scalar masses are not widely separated, and the charginos are not extravagantly heavy, will be accessible prior to LHC operation. The most sensitive tests for SUSY contributions in $B$ decays before the start of the LHC will come from $\\sin 2 \\beta$ measurements in charmonium modes and \\CPn -asymmetry measurements of $B \\rightarrow X_s \\gamma$ and $B \\rightarrow K^* \\gamma$ by BABAR and BELLE. For the latter studies the large SM theoretical uncertainties present in the branching-fraction measurements that make an extraction of SUSY contributions difficult are absent. In the LHC era improved measurements on $\\sin 2 \\beta$ and the \\CP asymmetry in $K^{*0} \\gamma$ will be carried out. In addition, measurements of branching-fractions and lepton forward-backward asymmetries of $B \\rightarrow s \\ell^+ \\ell^-$ modes will provide very sensitive tests for uncovering SUSY effects. While hadron-collider experiments will focus on $K^{(*)} \\mu^+ \\mu^-$ modes, asymmetric $B$-factory experiments will study $B \\rightarrow X_s \\ell^+ \\ell^-$ and $B \\rightarrow K^{(*)} \\ell^+ \\ell^-$ decays. At a super $B$ factory the precision can be significantly improved in these modes. While the first studies of $\\alpha$ will be performed by BABAR and BELLE, precise measurements will come from from BTEV and LHCb. In particular, the hadron-collider experiments produce high statistics $B_s$ samples, which allow for precise measurements of $\\gamma$. Additional precise measurements of $\\gamma$ at the $\\Upsilon(4S)$ can only come from a super $B$ factory. By this time a considerable reduction of the theoretical uncertainties in the extraction of $|V_{ub}/V_{cb}|$, $|V_{td}|$ and $|V_{ts}|$ from data is also expected, allowing to perform a sensitive model-independent analysis for extracting a new weak phase and new contributions to $B^0 \\bar B^0$ mixing. In order to observe the modes $ B \\rightarrow X_s \\nu \\bar \\nu$, $ B \\rightarrow K^+ \\nu \\bar \\nu$, or $ B \\rightarrow K^{*0} \\nu \\bar \\nu$, which bear the least theoretical uncertainties of radiative penguin decays, and study their properties a super $B$ factory is a prerequisite. So, for at least the next ten years a vivacious $B$ physics program ensures many high-precision measurements of \\CP asymmetries and radiative penguin decay properties that are suitable for extracting SUSY contributions." + }, + "0112/astro-ph0112065_arXiv.txt": { + "abstract": " ", + "introduction": "The GAIA mission will offer the wonderful opportunity to establish galactic dynamical models based on the data of a large number of stars. Indeed, complete positional and kinematical data over a large volume in the Galaxy are needed to construct well constrained models: we do not have such data for the moment, but this is precisely what the GAIA satellite will provide. We present here new tools to establish axisymmetric equilibrium models of the Milky Way, based on GAIA-like data. We know that our Galaxy is a barred one and so, the axisymmetric hypothesis is not correct, but axisymmetric models are a prerequisite for perturbation theory (for example the theory of spiral density waves) and thus they are still very useful. The models we wish to establish are pairs $(V,F)$ where $V$ is the gravitational potential generated by the whole mass distribution including the dark matter, and $F$ is the distribution function in phase space for late-type tracer stars in the galactic disk. For an equilibrium model (stationary potential and stationary distribution function solution of the collisionless Boltzmann equation), we know that the distribution function in phase space depends only on the integrals of the motion. During the last decade, many studies (e.g. \\cite{Du96}) have shown that the distribution function of tracer stars in the Milky Way has to depend on three isolating integrals of the motion, especially if information on the vertical motion of the stars is present as it will be with GAIA. In order to have an analytic third integral, in addition to the binding energy and the vertical component of the angular momentum, we use St\\\"ackel potentials \\cite{Z85}. ", + "conclusions": "" + }, + "0112/astro-ph0112253_arXiv.txt": { + "abstract": "Air showers recorded by the Haverah Park Array during the years 1974--1987 have been re-analysed. For the original estimate of the energy spectrum, a relationship between the ground parameter $\\rho$(600) and the primary energy, as determined by Hillas in the 1970s, was used. Here we describe the energy spectrum obtained using the QGSJET98 interaction model in the CORSIKA Monte Carlo code, together with GEANT to simulate the detailed detector response to ground particles. A new energy spectrum in the range 3 $\\times$ 10$^{17}$ eV to 4 $\\times$ 10$^{18}$ eV is presented. ", + "introduction": "\\label{intro.sec} After many years of research the origin of cosmic rays remains unclear. The composition of cosmic rays at high energies is not well known and even the energy spectrum is rather uncertain. Progress is limited primarily because cosmic rays above 10$^{15}$ eV can only be observed, with reasonable statistics, through the extensive air showers of secondary particles that they produce in the atmosphere. The primary energy and mass of the initiating cosmic ray have to be deduced from the form and particle content of the showers. This, unfortunately, requires use of an air shower model, based on our knowledge of electromagnetic and hadronic particle interactions and particle decay and transport in the atmosphere. These models suffer from poor knowledge of the hadronic and nuclear interactions at high energies and so the results inferred from the measurements are, to some extent, model dependent. At Haverah Park, UK, a 12 km$^2$ air shower array of water-Cherenkov detectors, was operational from 1967--1987 to measure cosmic rays in the energy range $10^{17}$ eV to $10^{20}$ eV. Here we present a re-analysis of data taken during 1974--1987. At that time the energy reconstruction of individual air shower events and the energy spectrum were obtained with a shower model of Hillas \\cite{hillas}, that gave an empirical description of high-energy interactions. In recent years a variety of more sophisticated models for hadron production, based on Gribov-Regge theory, have become available which attempt to include quantum chromodynamics in a consistent way. The QJSJET98 model (Quark-Gluon-String model with jets) \\cite{qgsjet} for high energy interactions, often used within the CORSIKA code \\cite{corsika} developed for Monte Carlo calculations of air showers, has been rather successful in describing consistently a variety of experimental results over a wide energy range. Tools for detailed simulation of the response of the detectors to particles of different type, energy and impact angle have also become available, e.g. GEANT \\cite{geant}, that allow us to study details of the measuring process in a manner inconceivable 25 years ago. The aim of this paper is to present a re-analysis of Haverah Park data on the energy spectrum using the QGSJET98 model in the CORSIKA code and with GEANT simulations. In the following paper, in this volume, we present also new results on the cosmic ray mass composition \\cite{masspaper}. ", + "conclusions": "The energy spectrum obtained in this analysis shows differences of up 30\\% at a given energy with a recent spectrum \\cite{watson} derived from Akeno and Haverah Park data. Our results are in good agreement with the recent results from the HIRes-MIA experiment \\cite{hires} in which a pure iron composition was assumed and with results reported at the Hamburg conference using monocular HiRes data \\cite{hiresdata}. The spectral index after the $knee$ at 3 $\\times$ 10$^{15}$ eV is estimated to be 3.1$\\pm$ 0.02 \\cite{kaskade}, so a further steepening of the primary energy spectrum between the knee and $3 \\times 10^{17}$ eV is needed to explain the spectral index derived in this work. It was claimed in \\cite{hires} that a break in the spectrum occurs at an energy $\\approx 3 \\times 10^{17}$ eV. Unfortunately this is the energy threshold of the spectrum described here. We have re-calculated the cosmic ray flux corresponding to the 4 highest energy events obtaining: JE$^3$=$8.27^{+6.5}_{-4.0}$ $\\times$ $10^{24}$ eV$^2$~sr$^{-1}$~m$^{-2}$~s$^{-1}$ at an energy of 7.62 $\\times$ $10^{19}$ eV. The average energy of these four events that were before above 10$^{20}$ eV, is shifted a $\\approx$ 30\\% downwards to energies below 10$^{20}$ eV. The most energetic event has an energy of 8.3 $\\times$ 10$^{19}$ eV. The Haverah Park array can be considered as an early prototype of the Auger Observatory which will employ water Cherenkov tanks of identical depth. The steps taken to derive an energy spectrum for the Auger Observatory are likely to be similar to those described here, and similar problems, like the attenuation length to be used and the $E$--$\\rho$(1000) relation, will need to be addressed." + }, + "0112/astro-ph0112123_arXiv.txt": { + "abstract": "s{Gravitational Lensing is a {\\sl unique} tool to constrain the mass distribution of collapsed structures, this is particularly true for galaxies, either on a case by case basis using multiple images of background sources (such as quasars), or statistically using the so called {\\it galaxy-galaxy lensing} technique. First, I will present the lensing theory, and then discuss the various methods applied to current observations. Finally, I will review the bright future prospects of galaxy lensing that will benefit of the development of high resolution, large, wide and deep (lensing) surveys. } ", + "introduction": " ", + "conclusions": "" + }, + "0112/astro-ph0112409_arXiv.txt": { + "abstract": "We describe the first results from a comprehensive study of the distant cluster Cl0024+1654 ($z=0.39$) based upon a pattern of 38 mosaiced HST-WFPC2 images extending to radii $\\sim 5$ Mpc. These are being analysed in conjunction with extensive spectroscopy conducted with the CFHT, WHT, and Keck Telescopes. The overall goal is to understand the morphological transformations and associated short-term star formation histories of representative numbers of infalling field galaxies in the context of the cluster potential as defined by weak lensing studies. Our HST database contains over 2000 galaxy morphologies to $I=22.5$. Spectroscopic data and HST morphologies are currently available for about 215 members over an unprecedented range of environments. We confirm the existence of a well-defined morphology-density relation over a large dynamic range within a single system at a significant look-back time. Tentative trends in the E/S0 fraction as a function of radius are discussed. A weak lensing signal in the background galaxies has been detected at the cluster periphery and its inversion demonstrates only marginal substructure. A statistically-significant galaxy-galaxy lensing signal has also been seen for the cluster members. Further work will relate radial dependencies in the dark matter and halo masses in the context of spectroscopic and morphological diagnostics of truncated star formation. ", + "introduction": "The ability of the Hubble Space Telescope (HST) to resolve galaxy morphologies at cosmologically-significant look-back times has enabled great progress in understanding the origin of the morphology density relation and, {\\em inter alia}, the influence of the environment on galaxy morphology. On the basis of HST images of several clusters at $z\\sim0.3-0.5$, the Morphs team (Dressler et al.\\ 1997) determined a remarkable decline with redshift in the fraction of S0s with a corresponding rise in that for spirals. They deduced that spirals were transformed to S0 galaxies remarkably recently, most likely because their gas was removed by processes such as tidal effects and ram pressure stripping. Alongside this evolutionary signal, strong evolution is also found in star formation and morphological characteristics of field galaxies (Glazebrook et al.\\ 1995; Lilly et al.\\ 1996). Understanding the role that gas-rich field galaxies play in fueling the transformations inferred in cluster cores remains an interesting issue. As field galaxies most likely fall into clusters at all epochs, the origin of the recent demise of spirals is puzzling. A promising way forward is to study, in considerable detail, the transformations occurring {\\em in situ} in a single system from the virialised core to the periphery where field galaxies infall on radial orbits. This is preferable to drawing deductions based on studies of many clusters, observed at different redshifts, each at an unknown stage in its evolutionary history. Because of the small field of view of WFPC-2, most morphological studies have been limited to the very central regions of distant clusters (e.g. Smail et al.\\ 1997), and correspondingly little is known about the properties of the infalling galaxy population at large radii. Ground based studies (Abraham et al.\\ 1996a) have explored the {\\em integrated} properties of galaxies to large cluster radii, finding evidence for radial trends in diagnostics of recent star formation. It seems likely that star formation is suppressed as galaxies infall, although the physical mechanisms and timescales are still quite controversial (Poggianti et al.\\ 1999; Balogh et al.\\ 2000). These latter uncertainties may be overcome if the timescales of stellar activity/truncation can be linked to those determined dynamically in a known gravitational potential (e.~g. from weak lensing signal). We present the first results from a wide field HST survey of the rich cluster Cl0024+1654 (hereafter 0024; $z=0.39$) whose goal is to address many of the above questions. In the following, the Hubble constant is H$_0=50h_{50}$kms$^{-1}$ Mpc$^{-1}$, assuming $h_{50}=1.3$ when necessary. The matter density of the Universe and the cosmological constant are $\\Omega=0.3$ and $\\Omega_{\\Lambda}=0.7$ respectively. ", + "conclusions": "" + }, + "0112/astro-ph0112315_arXiv.txt": { + "abstract": "{ We present a newly measured X-ray temperature function of galaxy clusters using a complete flux-limited sample of 61 clusters. The sample is constructed with the total survey area of 8.14 steradians and the flux limit of $1.99\\times 10^{-11}$ ergs s$^{-1}$ cm$^{-2}$ in the 0.1--2.4~keV band. X-ray temperatures and fluxes of the sample clusters were accurately measured with {\\it ASCA} and {\\it ROSAT} data. The derived temperature function covers an unprecedentedly wide temperature range of 1.4-11~keV. By fitting these data with theoretically predicted temperature functions given by the Press-Schechter formalism together with a recent formation approximation and the CDM power spectrum, we obtained tight and individual constraints on $\\Omega_{\\rm m,0}$ and $\\sigma_8$. We also employed the Formation-Epoch model in which the distribution in the formation epoch of clusters as well as the temperature evolution are taken into account, showing significantly different results. Systematics caused by the uncertainty in the mass-temperature relation are studied and found to be as large as the statistical errors. ", + "introduction": "The mass function of clusters of galaxies (MF), the number density of the most massive virialized systems, contains information on the structure formation history of the universe. A theoretical framework, e.g. the Press-Schechter formalism (Press \\& Schechter \\cite{PS}) together with the Cold Dark Matter model, has been established to predict the MF. This allows us to constrain cosmological parameters using an observationally determined MF for the present epoch (as well as its time evolution for even tighter constraints). In particular, $\\sigma_8$, the amplitude of mass density fluctuations on a scale of 8$h^{-1}$ Mpc where $h$ is the Hubble constant in unit of 100 km/s/Mpc, and $\\Omega_{\\rm m,0}$, the mean matter density, are most sensitively determined by the cluster abundance measurements (e.g. Henry \\& Arnaud \\cite{HA}; White et al. \\cite{WEF}; Eke et al. \\cite{ECF}; Kitayama \\& Suto \\cite{KS}; Viana \\& Liddle \\cite{VL}; Oukbir \\& Blanchard \\cite{OB}; Pen \\cite{Pen}; Eke et al. \\cite{Eke}). Observationally the local MF has been derived from measuring masses of individual clusters from galaxy velocity dispersions or other optical properties by Bahcall and Cen (\\cite{BC}), Biviano et al. (\\cite{Biviano}), and Girardi et al. (\\cite{Girardi}). The estimated virial masses for individual clusters depend rather strongly on model assumptions, however. As argued by Evrard et al. (\\cite{Evrard97}) on the basis of hydrodynamical N-body simulations, cluster masses may be presently more accurately determined from a temperature measurement and a mass-temperature relation determined from detailed observations or numerical modeling. Thus alternatively, as a well-defined observational quantity, the X-ray temperature function (XTF) has been measured, which can be converted to the MF by means of the mass-temperature relation. The first measurements of the XTF were reported by Edge et al. (\\cite{Edge}) and Henry \\& Arnaud (\\cite{HA}), using an X-ray flux-limited sample of 45 clusters and 25 clusters, respectively. Recent observational improvement on the XTF was made by Markevitch (\\cite{Markevitch}), Henry (\\cite{Henry}), Blanchard et al. (\\cite{Blanchard}), and Pierpaoli et al. (\\cite{Pierpaoli}), using more accurate temperature-measurement results for each cluster with {\\it ASCA} data (Tanaka et al. \\cite{ASCA}). However, the narrow temperature ranges (3-10~keV) in which the XTF is defined so far do not allow an investigation of a more detailed shape of the XTF other than just fitting a single power-law. Without better information on the actual shape of the XTF, no independent constraints on $\\sigma_8$ and $\\Omega_{\\rm m,0}$ can be derived, and a wide range of combinations is still allowed. Recently, a number of new X-ray-cluster surveys were performed, which provide a high completeness for the brightest clusters. This motivated us to revisit the measurement of the local XTF and to improve its accuracy in order to derive narrower constraints on the cosmological parameters. Reiprich \\& B\\\"{o}hringer (\\cite{RB}) have compiled a new X-ray flux-limited cluster sample ({\\sfsl HIFLUGCS}) with a flux limit of $2\\times10^{-11}$ ergs s$^{-1}$ cm$^{-2}$ (0.1-2.4~keV). Compared with previous samples with similar flux limits, their catalog covers the largest volume and is the most complete. Based on this cluster sample, Reiprich \\& B\\\"{o}hringer have measured total masses for individual clusters and derived the X-ray mass function for the first time. In this paper we, using this cluster sample, report the construction of a new measurement of the XTF. The larger number of clusters in the sample together with accurate temperature measurements with {\\it ASCA} data leads to a significant improvement of the XTF, which covers 1.4-11~keV temperature range. The X-ray flux-limited sample and the new temperature measurements with {\\it ASCA} data are presented in Sect. 2. The derivation of the XTF is described in Sect. 3. In Sect. 4, the XTF is used to obtain constraints on cosmological parameters, and the results are compared with other work in Sect. 5. Throughout the paper, the Hubble constant is given as $100\\ h$ km/s/Mpc, and $\\log$ and $\\ln$ denotes a decimal and natural logarithm, respectively. ", + "conclusions": "The most important result concerning the XTF derived here is that it covers a large enough range of temperatures to reveal the shape of the function beyond a simple power law representation. As expected from theoretical considerations, there should be a very sharp cut-off at the high mass end and the function should turn into a shallower slope at the low mass end as shown by our observations. We have demonstrated that these constraints on the actual slope of the XTF allow us to derive independent constraints on the two parameters which are most important in cosmic structure formation models: $\\Omega_{\\rm m,0}$ and $\\sigma_8$. Leaving out only 16\\% of the sample clusters from the low temperature end almost doubles the uncertainty for these independent constraints. Therefore, while these independent constraints are an important achievement of the present work on one hand, this demonstration also shows that these constraints are sensitive to any systematic error introduced into the XTF. For this reason we have studied the stability of our results to variations introduced to the observed data and also variations of the theoretical modeling. We have fully incorporated the measurement errors in fluxes and temperatures of the sample clusters. With the typical errors of 5\\% for the flux and 6-10\\% for the temperature, the biases were found to be relatively small and not to cause a drastic change on the final results in our analysis. On the other hand, we have shown that there are considerable variations and uncertainties in the theoretical modeling of the XTF, which introduce significant systematic errors in the final constraints on $\\Omega_{\\rm m,0}$ and $\\sigma_8$. Various $M-T$ relations that are employed to model the XTF have resulted in considerably different values of $\\Omega_{\\rm m,0}$ and $\\sigma_8$. Another source of variation in the theoretical modeling that we have studied is a recent formation approximation which is employed with the Press-Schechter formalism to build the PS model here. Using the FE model, in which the distribution of cluster formation redshifts and the evolution of the mass as well as temperature of a cluster after the collapse are analytically formulated, we have shown that a considerable difference could be introduced in the derived cosmological parameters, in particular in $\\sigma_8$. This further motivates us to establish a more sophisticated model to the XTF incorporating new X-ray observations of the high redshift clusters with Chandra and XMM-Newton. Finally, taking into account all the uncertainties, we put our general constraints on $\\Omega_{\\rm m,0}$ and $\\sigma_8$ as $\\Omega_{\\rm m,0}$=0.03--0.37, $\\sigma_8$=0.52--1.15 in an open universe and $\\Omega_{\\rm m,0}$=0.06--0.38, $\\sigma_8$=0.57--1.33 in a flat universe, respectively. From the comparison with the {\\it COBE} constraints, a flat universe is more preferable than an open universe. \\appendix" + }, + "0112/astro-ph0112159_arXiv.txt": { + "abstract": " ", + "introduction": "\"Old\" questions of relativistic astrophysics, like the internal structure of compact stars, have received a lot of attention recently since they are very important for a solid knowledge of the QCD diagram in the low-$T$ , high $\\mu$ region as discussed along this Workshop. Among the most simple forms of investigating (indirectly) the nature of high-density matter, more precisely of the equation of state connecting the pressure and the energy density; is by obtaining accurate determinations of masses and radii. A comparison of the static models generated by integration of the Tolman-Oppenheimer-Volkoff equations with observed data should reveal a great deal of information about the equation of state. While astronomers and physicists alike have longely dreamed of such determinations we have only recently achieved sufficient accuracy to perform some key tests on selected objects, although it should be acknowledged that three decades of compact star astrophysics had produced quite clever arguments to determine masses and radii, even in the cases in which they proved to be wrong. A paradigmatic example of the above assertions is the celebrated determination of the binary pulsar mass PSR 1913+16, believed to be accurate to the fourth decimal place \\cite{vankerkwijk}. Methods based on combinations of spectroscopic and photometric techniques have been recently perfectioned, and have confirmed that at least one X-ray source is significantly above the \"canonical\" 1.44 $M_{\\odot}$; namely Vela X-1 for which a value of $1.87^{+0.23}_{-0.17} \\, M_{\\odot}$ has been obtained \\cite{vankerkwijk}. It is also possible that at least some compact stars are extremely compact, or in other words, that their radii are $\\sim 30-40 \\%$ less than the cherished $10 \\, km$ for $M \\, \\sim 1 \\, M_{\\odot}$. Indeed, this is the claim of the analysis of \\cite{li} of the binary Her X-1 ($M = 0.98 \\pm 0.12 \\, M_{\\odot}$ and $R = 6.7 \\pm 1.2 \\, km$) and \\cite{pons} in the case of the isolated nearby RX J1856-37 ($M = 0.9 \\pm 0.2 \\; M_{\\odot}$ and $R = 6^{-1}_{+2} \\, km$). In fact, such compactness is extremely difficult (impossible?) to model using underlying equations of state based on hadrons alone, and a \"natural\" alternative would be to consider deconfined matter. Naturally, these results will be checked and scrutinized for confirmation. Could this be related to high-density deconfined matter ? Perhaps, since there is hope to achieve deconfinement densities inside compact stars. Also, there is some evidence that diquarks (e. g. a spin-0, color-antitriplet bound state of two quarks) might occur as a component in the QCD plasma (see \\cite{anselmino} and references therein for a review). Such diquarks would be expected to be favoured by Bose statistics and they are quit helpful to model the low-energy hadronic properties. Of course, at very high densities, characterized by interquark distances less than $(10 GeV^2 )^{-1/2}$, diquarks lose their identity and must eventually dissolve into quarks, even if there is no clear consensus about the onset of the asymptotic regime. According to this picture we may regard a diquark as any system of two quarks considered collectively. Diquark correlations arise in part from spin-dependent interactions between two quarks. Regardless of the exact mechanism for their origin, it is imperative to have some gain of energy (binding) for the stellar models to work (see below). The bound state is a quite strong assumption, and is not easy to prove (or disprove) in the absence of a reliable way of performing sensible calculations. As a working hypothesis (quite analogous to the well known strange matter), we shall assume simply that this bound state actually exists. ", + "conclusions": "Three examples of the QMDD diquark model are shown in Fig. 1. The resulting equations of state are very soft (models B and C), as expected, but fall short to explain the claimed compactness by at least $\\Delta R \\sim 1 \\, km$, even in the most favourable case (model B). When compared with the equation of state for quark-diquark matter presented in \\cite{horvath1} (model A) and the softest models of MIT bag models \\cite{MIT} (model D) we may assert that the stellar sequences are in between the former and the latter. Diquarks may be relevant for compact star structure because they allow the existence of stable compact stars with masses of $\\sim 1 M_{\\odot}$ and very small radii. For comparison, the inferred radius and mass of both Her X-1 and RX J1856-37 are shown with their respective error bars. Li {\\it et al.} models SS1 and SS2 \\cite{ignazio}(not shown in the figure) are examples of other self-bound models of the equation of state which do produce very compact stars, in fact also capable of matching the values of Her X-1 claimed in \\cite{li} and of RX J1856-37 \\cite{pons}. We may state that if evidence for extreme compactness holds, diquarks may be relevant for model building . \\begin{figure}[htb] \\begin{center} \\epsfig{file=bohr.eps,height=2.5in} \\caption{The mass-radius plane, stellar sequences labeled as indicated in the text} \\label{fig:bohr} \\end{center} \\end{figure} \\bigskip G. Lugones acknowledges the IAG-S\\~ao Paulo for hospitality and the financial support received from the Funda\\c c\\~ao de Amparo \\`a Pesquisa do Estado de S\\~ao Paulo. J.E. Horvath wishes to acknowledge the CNPq Agency (Brazil) for partial financial support. The Scientific Organizers of the meeting and Mrs. Ana Rey are acknowledged for making JEH stay in Copenhaguen a very pleasant experience and partial financial support to attend the event. I. Bombaci, S. Balberg and S. Fredricksson have shared with him several interesting discussions." + }, + "0112/astro-ph0112190_arXiv.txt": { + "abstract": "We present in catalog form the optical identifications for objects from the first phase of the Wide Angle ROSAT Pointed Survey (WARPS). WARPS is a serendipitous survey of relatively deep, pointed ROSAT observations for clusters of galaxies. The X-ray source detection algorithm used by WARPS is Voronoi Tessellation and Percolation (VTP), a technique which is equally sensitive to point sources and extended sources of low surface brightness. WARPS-I is based on the central regions of 86 ROSAT PSPC fields, covering an area of 16.2 square degrees. We describe here the X-ray source screening and optical identification process for WARPS-I, which yielded 34 clusters at $0.06 1$ for clusters). WARPS uses the Voronoi Tessellation and Percolation (VTP, \\citealt{ebeling93,ew93}) algorithm to detect both extended and point-like X-ray sources. Our application of VTP to detect X-ray sources has been described in Paper I, as has the survey calibration. An extensive optical follow-up program (described in detail in Paper II) imaged not only extended X-ray sources (which are the most promising cluster candidates) but also all blank field point X-ray sources, and then took spectra of galaxies in the fields of cluster candidates. We assume that groups and clusters of galaxies form a continuous population, referred to simply as ``clusters,'' and do not further distinguish between groups and clusters of galaxies. Early measurements of the cluster XLF, based on EXOSAT data (\\citealt{Ed90}) and the {\\it Einstein} Extended Medium Sensitivity Survey (EMSS, \\citealt{He92}), indicated strong negative evolution at $L(0.3-3.5{\\rm ~keV}) >10^{44} {\\rm ~erg ~s^{-1}}$, perhaps even at relatively low redshifts. This was seen as support for critical density, cold dark matter cosmologies (e.g., \\citealt{vialit96}). However, a large number of recent results have changed the picture substantially. At low redshift ($z<0.3$), the results of the ROSAT Brightest Cluster Sample (BCS, \\citealt{bcsII}) show little if any evolution of the cluster XLF. At higher redshift, both new ROSAT surveys (RDCS, \\citealt{Ros98}; WARPS, Papers II, III; SHARC-south \\citealt{Bur97}; SHARC-Bright \\citealt{sharc2000}; CfA 160 deg$^2$ [also known as VMF], \\citealt{vmf98}; NEP, \\citealt{mu00}) and re-analyses of EMSS data (\\citealt{Nic97}, Stocke et al., in preparation) are again consistent with no evolution at $z<1$ (except perhaps at the highest luminosities; see in particular \\citealt{Ros98,vmfII}). The reason for this discrepancy is a subject of active debate, but contributing factors may include underestimation of X-ray flux by the EMSS (Papers II, III), and misidentification of some EMSS X-ray sources as clusters (\\citealt{Nic97, Rec99}). Indeed, the most current analysis of the highest-redshift part of the EMSS sample (\\citealt{LG95}) is also consistent with no significant evolution between $z\\sim 0.33$ and $z\\sim 0.75$ (see also Paper III; although for a somewhat different view of the same literature see \\citealt{nep}). Here we describe the identification process and and present in catalog form optical identifications for objects from the first phase of the WARPS project (WARPS-I), which included data from 86 ROSAT Position Sensitive Proportional Counter (PSPC) fields. This sample formed the basis of Papers I and II. The optical identifications for galaxy clusters from WARPS-I are essentially complete. Papers III, IV, and V also include data from the second phase of the WARPS project (WARPS-II), which includes nearly 4 times as many PSPC fields and concentrates on the most distant clusters. Those identifications are almost complete and will be presented elsewhere (Horner et al. in preparation). In \\S~\\ref{sec:methods} we describe briefly the survey methods and source identification process, as well as cross-correlations with various astronomical databases. In \\S~\\ref{sec:efficacy} we comment explicitly on the efficacy of VTP as a method for detecting clusters of galaxies, and compare it with other X-ray source detection methods currently in use. In \\S~\\ref{sec:clusters} we describe the statistically complete, flux limited WARPS-I sample of clusters of galaxies, and also comment on individual clusters and interesting sources which fall below the flux limit. We conclude in \\S~\\ref{sec:prospects} by outlining future directions for research with this and other cluster samples. Unless otherwise stated, we use $q_{0} = 0.5$ and $H_{0} = 50$ km s$^{-1}$ Mpc$^{-1}$ when calculating distance dependent quantities. Similarly, we use the 0.5--2.0 keV ROSAT PSPC band when quoting count rates, fluxes, and luminosities unless otherwise noted. The catalogs, finder charts, and other information can be found on our WARPS WWW page at {\\tt \\url{http://lheawww.gsfc.nasa.gov/\\~\\ horner/warpsI/warpsI.html}}. ", + "conclusions": "\\label{sec:prospects} Here we have presented the identifications and cluster catalogs from the first phase of the WARPS project, an extensive survey of deep ROSAT pointings using a surface brightness sensitive algorithm. In doing so, we have also presented some comparisons of the results of WARPS-I to other cluster surveys, which help place the the various cluster detection algorithms and follow-up programs in perspective. As can be seen, the findings of WARPS are largely similar to those of other cluster surveys as regards the space density of X-ray emitting clusters. The number of serendipitous X-ray surveys for clusters of galaxies now number more than a dozen. For example, the list of those based on ROSAT data includes BCS \\citep{bcsI,bcsII,bcsIII}, WARPS (Papers I, II), RDCS \\citep{Ros98}, SHARC-Bright \\citep{sharc2000}, SHARC-south \\citep{Bur97,Col97}, CfA 160 deg$^2$ \\citep{vmf98}, NEP \\citep{nep}, NORAS \\citep{noras}, REFLEX \\citep{reflexI}, MACS \\citep{macs} and CIZA \\citep{ciza}. Each of these surveys use different selection criteria and/or X-ray source-detection algorithms, but largely similar optical follow-up programs. Meanwhile, several optical surveys have used moderate depth, large-area images of the sky combined with algorithms that measure overdensities of galaxies (e.g., PDCS, \\citealt{pdcsI, pdcsII}; RCS, \\citealt{glye}). All of these surveys represent a fertile ground for constraining cosmological parameters as well as learning more about appropriate methods for finding massive clusters of galaxies and their attendant selection effects. However, precious few efforts (most notably ROXS, \\citealt{roxs}) are underway to understand the relationship of the selection effects in optical and X-ray surveys. Despite the different emphases inherent in each method, all of the new surveys show no evidence for evolution at all but the most X-ray luminous clusters of galaxies at redshifts up to $z=0.8$. The X-ray luminosity-temperature relationship (Paper IV, \\citealt{MS97}) at $z<0.8$, also shows no sign of evolution. Yet at $z \\sim 1$, there are tantalizing hints of evolution in the cluster population. This evidence includes a preponderance of irregular X-ray morphologies, optical cluster galaxy distributions and colors (Paper III, \\citealt{sta01}, although see Paper V for a notable counter-example), as well as a suggestion that the space density of the most luminous ($L_X > 5 \\times 10^{44} {\\rm ~erg ~s^{-1}}$) clusters decreases significantly at redshifts $z>0.5$ (\\citealt{Ros98,vmfII}; although for a counterargument see Paper V). Taken together, these lines of evidence might suggest that the epoch of cluster formation is currently just out of reach of our current observational tools, which are limited to finding clusters at $z<1.25$ by the limitations inherent in observing in the optical (at $z=1.25$ the Ca break passes out of the I band), even though simulations ({\\it e.g.}, Paper I) show an ability to detect X-ray luminous, extended sources at redshifts up to nearly $z=2$. This presents future X-ray surveys ({\\it e.g.}, {\\it Chandra} or XMM-Newton) with significant challenges, since one would expect them to contain large numbers of $z>1$ clusters. Indeed, as we pointed out in Papers III and V, WARPS is sensitive enough to include clusters of $L_X \\sim 10^{45} {\\rm ~erg ~s^{-1}}$ out to redshifts of nearly 2. Even given the highly unclear (and very poorly explored) state of galaxy evolution at $z\\gtrsim 1$, it is obvious from the current literature that, at $z\\sim 0.8$, the space density and ICM characteristics of clusters of galaxies show little difference from those observed at lower redshifts. It is therefore imperative that future surveys for clusters of galaxies include near-infrared observations, which should increase their window of sensitivity for identifying cluster counterparts to X-ray sources out to $z\\sim 4.5$ (i.e., when the H \\& K break passes out of the K band). Of the current X-ray surveys, only RDCS has included near-infrared observations in its follow-up program, and as a result, RDCS is the only X-ray selected sample of clusters to contain $z>1$ objects \\citep{ros99, ros00}, albeit only a handful (4 according to \\citealt{ros00} and P. Rosati, private communication from an incomplete survey of about 80\\% of RDCS blank-field X-ray sources). The now overwhelming evidence of little or no significant evolution in the space density of clusters at redshifts up to 0.8 is difficult to reconcile with $\\Omega_m=1$ cold dark matter cosmologies, which predict significant evolution in the cluster LF at $z\\gtrsim 0.3$ . However, as \\citet{bod01} show, a great deal of degeneracy remains in the predictions of hierarchical models of structure formation at $z>1$, and coincidentally, this is where the fewest such X-ray selected clusters are known. It is at these highest redshifts that future X-ray and other surveys for clusters (notably Sunyaev-Zel'dovich effect methods, see \\citealt{xuwu01}) will have to concentrate their efforts and where they will have the most impact." + }, + "0112/astro-ph0112473_arXiv.txt": { + "abstract": "\\vspace{1pc} We report the analysis of the N$_e$-N$_\\mu$ coincident events collected by the MACRO/EAS-TOP collaboration at the Gran Sasso Laboratories. The result points to a primary composition becoming heavier around the knee of the primary spectrum (in the energy region $10^{15}-10^{16}$ eV). The result is in very good agreement with the measurements of EAS-TOP alone at the surface, wich detects muons with energy $E_{\\mu} > 1$ GeV and uses the same (QGSJET) interaction model. ", + "introduction": "The study of the primary composition in the Extensive Air Shower energy region requires the use of different observables in order to cross check the information and reduce the dependence on the interaction model and propagation codes used. At the National Gran Sasso Laboratories a program of exploiting the surface shower size measurements from EAS-TOP (2005 m a.s.l.), and the high energy muon measurements ($E_{\\mu}^{th} = 1.3 $ TeV) performed in the deep underground laboratories (MACRO) has been developed. Such muons in fact originate from the decays of mesons produced in the first interactions in the atmosphere and from a quite different rapidity region than the GeV muons usually used for such analysis ($x_{F}> 0.1$ or $ 0.2$). The two experiments operated in coincidence for a live time of $\\Delta T =$ 960.12 days between 1990 and 2000. The number of coincident events collected with the two detectors operating in the final configuration amounts to 28160, of which 3752 have shower cores inside the edges of the array (\"internal events\") and shower size $N_e > 2.10^5$, and 409 have $N_e > 10^{5.92}$, i.e. above the observed knee at the corresponding zenith angle. We present here an analysis of the full data set in terms of the QGSJET interaction model as implemented in CORSIKA. For a comparison of different interaction models see \\cite{as1}. The EAS-TOP array is located at Campo Imperatore (2005 m a.s.l., $ \\approx 30^o$ with respect to the vertical of the underground Gran Sasso Laboratories, corresponding to 930 gr cm$^{-2}$ atmospheric depth). Its e.m. detector (to which we are mainly interested in the present analysis) is built of 35 scintillator modules 10 m$^2$ each, including an area A $\\approx$ 10$^5$ m$^2$. The array is fully efficient for $N_e > 10^5$. Its reconstruction capabilities of the EAS parameters (for internal events) are: ${{\\Delta N_e} \\over N_e} \\approx 10 \\%$ above $N_e \\approx 10^5$, and $\\Delta \\theta \\approx 0.9^o$ for the EAS arrival direction. The array and the reconstruction procedures are fully described in \\cite{as2}. MACRO, in the underground Gran Sasso Laboratory at 963 m a.s.l., with 3100 m w.e. of minimum rock overburden, is a large area multi-purpose apparatus designed to detect penetrating cosmic radiation. The lower part of the MACRO detector has dimensions $76.6 \\times 12 \\times 4.8$ m$^3$. A detailed description of the apparatus can be found in \\cite{as21}. In this work we consider only muon tracks which have at least 4 aligned hits in both views of the horizontal streamer tube planes over the 10 layers composing the whole detector. The standard reconstruction procedure of MACRO \\cite{as3} has been used. The two experiments are separated by a thickness of rock ranging from 1100 up to 1300 m, depending on the angle. The muon energy threshold at the surface for muons reaching the MACRO depth ranges from $E_{\\mu}^{th} = 1.3$ TeV to $E_{\\mu}^{th} = 1.8$ TeV inside the effective area of EAS-TOP. Event coincidence is established off-line, using the absolute time given by a GPS system with an accuracy better than 1 $\\mu s$. Independent analyses of the two arrays are in \\cite{as4} and \\cite{as5}. ", + "conclusions": "The analysis of the N$_e$-N$_\\mu$ ($E_{\\mu} > 1.3$ TeV) data collected by the MACRO/EAS-TOP collaboration at the Gran Sasso Laboratories points to a primary composition becoming heavier around the knee of the primary spectrum (in the energy region $10^{15}-10^{16}$ eV). The result is in very good agreement with the measurements of EAS-TOP alone at the surface, wich detects muons with energy $E_{\\mu} > 1$ GeV, using the same (QGSJET) interaction model \\cite{as6}. GeV and TeV muons are produced in different kinematic regions: in the central one and at the edges of the fragmentation region respectively. The present measurements show therefore a good consistency of the interaction model in describing the yield of secondaries over a wide rapidity region." + }, + "0112/astro-ph0112535_arXiv.txt": { + "abstract": "{ We have reconstructed the galactic orbits of the parent stars of exoplanets. For comparison, we have recalculated the galactic orbits of stars from the \\cite{edva93} catalog. A comparison between the two samples indicates that stars with planets are not kinematically peculiar. At each perigalactic distance stars with planets have a metallicity systematically larger than the average for the comparison sample. We argue that this result favors scenarios where the presence of planets is the cause of the higher metallicity of stars with planets. ", + "introduction": "Spectroscopic analysis of parent stars of exoplanets had shown that these stars are more metal-rich than field stars (\\citealt{gon96,but00}). Two scenarios had been proposed to explain the high metallicity: (i) During the build-up of planets, the gravitational interaction among them (or with the disk) injects some objects in high-eccentricity orbits that can intersect the surface of the host star. If a sufficient number of these objects are captured by the star, then the photospheric metallicity will be enhanced by the dissolution of the planet. (ii) Planet formation is enhanced by the high metallicity of the parent protostar nebula. If there is a correlation between metal abundances and other properties of the stars that should have no impact on the formation of planets (e.g. their kinematics), we may expect systematic differences between stars with planets and without planets, in the first scenario, while there should be no difference in the second scenario. In this paper we study the second hypothesis. Following the work by \\cite{asa91}, using stellar parallaxes and proper motions from Hipparcos, we have reconstructed the galactic orbit of the parent stars of exoplanets. For comparison, we have recalculated the galactic orbit from the \\cite{edva93} catalogue, adopting a solar galactocentric distance of 8.5 kpc instead of the 8.0 kpc used by Edvardsson. A comparison between the two samples indicates that they are quite similar, and therefore that the stars with planets are not kinematically peculiar. ", + "conclusions": "Our data may be used to discuss the metallicity-planet connection. In fact, if the high metallicity is the cause for the presence of planets, we should not expect any correlation between presence of planets and galactocentric distance: it is the overall metallicity that is important, and in a given metallicity bin the distribution of stars with perigalactic distance should be the same for stars with and without planets. On the other hand, if planet capture is the mechanism that enhances the metallicity, we should expect that the SWP are the upper envelope of the distribution of metallicity with galactocentric distance: at any galactocentric distance the stars with planets should be more metal-rich than average, and in a given metallicity bin, stars with planets should have on average smaller perigalactic distances. At any galactocentric distance, the metallicity of the stars with planets is roughly the upper envelope of the metallicity distribution of the comparison sample. The few known (mildly) metal-poor stars with planets all have small values of the perigalactic distance, and there are no stars with known planets among the much more frequent (in the solar neighborhood) mildly metal-poor stars that have nearly circular orbits. Based on the relative frequency in the solar neighborhood of mildly metal-poor stars on circular and highly eccentric orbits, we should expect a large number of mildly metal-poor stars with planets. These are not observed. A Kolmogorov-Smirnov test for the mildly metal-poor ([Fe/H] $\\le -0.1$) SWP and Edvardsson stars gives a probability that both distributions with perigalactic distances were extracted from the same parent population of 0.0046; the same test for the higher value of metallicity ([Fe/H] $> -0.1$) gives a probability of 0.5159. This result does not depend critically on possible offsets between the abundance scales used for SWP and the Edvardsson sample: in fact, even if we lower [Fe/H] for all SWP by 0.07~dex (the systematic offset measured for the nine stars in common in the two samples), the possibility that the mildly metal-poor stars ([Fe/H] $\\le -0.1$) in the two samples were extracted from parent populations having the same distribution with perigalactic distance can still be rejected at a quite good level of confidence (the probability that they are drawn from the same distribution is 0.022). The test for the higher metallicity objects gives a probability of 0.6644. Note that this is not due to a selection effect related to the apparent brightness of the stars: in fact in the Edvardsson sample, metal-poor stars ([Fe/H] $\\le -0.1$) with $R_p \\le 7$ kpc are on average fainter than metal-poor stars with $R_p > 7$ kpc (average magnitudes of 6.16 $\\pm$ 0.20 and 5.80 $\\pm$ 0.12 respectively). It should then be easier to discover planets around the brighter stars with near circular orbits than around the fainter stars with eccentric orbits. In the scenario where high metallicity is the cause of the presence of planets, when the metallicity exceeds a critical value, the planets are present, independent of the perigalactic distance. Hence we expect that in a given metal abundance range, the number of stars with planets at different perigalactic distance will follow the distribution of the parent population of all nearby stars. We should then find more stars with planets with large perigalactic distance than with small perigalactic distance (since the parent population is more numerous). In the opposite scenario, where the planets are the cause of the high metallicity of the central stars, the SWP are the upper envelope of the distribution of [Fe/H] versus perigalactic distance of field stars: this is because at any perigalactic distance, stars with planets are metal enriched with respect to stars without planets. In this case we expect that in the mildly metal-poor abundance bin, we may find a larger number of stars with small perigalactic distance than stars with large perigalactic distance, because the metal enrichment caused by the presence of the planets will push the stars with planets extracted from the parent population of mildly metal-poor stars out of this metallicity bin (into the high metallicity bin). We conclude that insofar as no kinematic selection effect is present in the sample of stars with planets, the fact that planets have been found only among mildly metal-poor stars (presently in the solar neighborhood) that have small perigalactic radii (in spite of the relative rarity of such objects) clearly favours the second scenario. We argue that this result strongly favours scenarios where the presence of planets is the cause of the higher metallicities. Of course our results do not rule out the possibility that higher metallicity also favours the presence of planets (i.e. that both scenarios are applicable)." + }, + "0112/astro-ph0112229_arXiv.txt": { + "abstract": "In Paper III of our series $``$A Uniform Analysis of the Ly-$\\alpha$ forest at $z=0 - 5$\", we presented a set of 270 quasar spectra from the archives of the Faint Object Spectrograph (FOS) on the Hubble Space Telescope (HST). A total of 151 of these spectra, yielding 906 lines, are suitable for using the proximity effect signature to measure $J(\\nu_{0})$, the mean intensity of the hydrogen-ionizing background radiation field, at low redshift. Using a maximum likelihood technique and the best estimates possible for each QSO's Lyman limit flux and systemic redshift, we find $J(\\nu_{0})= 7.6^{+9.4}_{-3.0} \\times 10^{-23}$ ergs s$^{-1}$ cm$^{-2}$ Hz$^{-1}$ sr$^{-1}$ at $0.03 < z < 1.67$. This is in good agreement with the mean intensity expected from models of the background which incorporate only the known quasar population. When the sample is divided into two subsamples, consisting of lines with $z < 1$ and $z > 1$, the values of $J(\\nu_{0})$ found are 6.5$^{+38.}_{-1.6} \\times 10^{-23}$ ergs s$^{-1}$ cm$^{-2}$ Hz$^{-1}$ sr$^{-1}$, and 1.0$^{+3.8}_{-0.2} \\times 10^{-22}$ ergs s$^{-1}$ cm$^{-2}$ Hz$^{-1}$ sr$^{-1}$, respectively, indicating that the mean intensity of the background is evolving over the redshift range of this data set. Relaxing the assumption that the spectral shapes of the sample spectra and the background are identical, the best fit HI photoionization rates are found to be $6.7 \\times 10^{-13}$ s$^{-1}$ for all redshifts, and $1.9 \\times 10^{-13}$ s$^{-1}$ and $1.3 \\times 10^{-12}$ s$^{-1}$ for $z < 1$ and $z > 1$, respectively. The inclusion of blazars, associated absorbers, or damped Ly-$\\alpha$ absorbers, or the consideration of a $\\Lambda$ CDM cosmology rather than one in which $\\Omega_{\\Lambda}=0$ has no significant effect on the results. The result obtained using radio loud objects is not significantly different from that found using radio quiet objects. Allowing for a variable equivalent width threshold gives a consistently larger value of $J(\\nu_{0})$ than the constant threshold treatment, though this is found to be sensitive to the inclusion of a small number of weak lines near the QSO emission redshifts. This work confirms that the evolution of the number density of Ly-$\\alpha$ lines is driven by a decrease in the ionizing background from $z \\sim 2$ to $z \\sim 0$ as well as by the formation of structure in the intergalactic medium. ", + "introduction": "\\label{sec-intro} The spectra of quasars show a $``$forest\" of absorption lines blueward of the Ly-$\\alpha$ emission line (Lynds 1971, Sargent et al.\\ 1980, Weymann, Carswell, \\& Smith 1981). Observational and theoretical work in recent years has shown that most of this absorption can be attributed to neutral hydrogen in galaxies and large-scale structure along the line of sight (Cen et al.\\ 1994, Lanzetta et al.\\ 1995,1996, Zhang et al.\\ 1995, Hernquist et al.\\ 1996, Miralda-Escud\\'{e} et al.\\ 1996, Bi \\& Davidsen 1997, Chen et al.\\ 1998, Theuns et al.\\ 1998, Ortiz-Gil et al.\\ 1999, Impey, Petry, \\& Flint 1999, Dav\\'{e} et al.\\ 1999, Bryan et al.\\ 1999). In aggregate, QSO spectra show an increasing line density with increasing redshift such that $dN/dz \\propto (1.+z)^{\\gamma}$ (Sargent et al.\\ 1980, Weymann, Carswell, \\& Smith 1981, Young et al.\\ 1982, Murdoch et al.\\ 1986, Lu, Wolfe, \\& Turnshek, 1991, Bechtold 1994, Kim et al.\\ 1997). But the line density within an individual quasar spectrum decreases with proximity to the Ly-$\\alpha$ emission line (Weymann, Carswell, \\& Smith 1981, Murdoch et al.\\ 1986). This is generally thought to be due to enhanced ionization of neutral hydrogen in the vicinity of the quasar due to ionizing photons from the quasar itself. This $``$proximity effect\" can be used to measure the mean intensity of the UV background, denoted $J(\\nu_{0})$ (Carswell et al.\\ 1987, Bajtlik, Duncan, \\& Ostriker 1988, hereafter BDO). $J(\\nu_{0})$ has been measured at $z > 1.7$ by a variety of authors (BDO, Lu, Wolfe, \\& Turnshek 1991, Giallongo et al.\\ 1993,1996, Bechtold 1994, Williger et al.\\ 1994, Cristiani et al.\\ 1995, Fern\\'{a}ndez-Soto et al.\\ 1995, Lu et al.\\ 1996, Savaglio et al.\\ 1997, Cooke et al.\\ 1997, Scott et al.\\ 2000b). The results, summarized in Paper II of this series (Scott et al.\\ 2000b), are in general agreement with the predictions of models of the UV background which integrate the contribution from known population of quasars and include reprocessing effects in an inhomogeneous intergalactic medium (Haardt \\& Madau 1996, hereafter HM96, Fardal et al.\\ 1998). In Paper II, the mean intensity of the ionizing background was found to be $7.0^{+3.4}_{-4.4} \\times 10^{-22}$ ergs s$^{-1}$ cm$^{-2}$ Hz$^{-1}$ sr$^{-1}$ at $z \\sim 3$. The decline of the quasar space density from $z \\sim 2$ to the present is expected to drive a corresponding decline in the intensity of the UV background. Kulkarni \\& Fall (1993, hereafter KF93) measured $J(\\nu_{0}) \\sim 6 \\times 10^{-24}$ ergs s$^{-1}$ cm$^{-2}$ Hz$^{-1}$ sr$^{-1}$ at $z \\sim 0.5$ from a subset of the now complete HST Quasar Absorption Line Key Project sample presented by Bahcall et al.\\ (1993). Much of this previous work has relied upon the technique for measuring $J(\\nu_{0})$ outlined by BDO. This technique requires the entire sample of absorption lines to be binned according to the ratio of the quasar flux at the physical position of the absorber to the background flux. This is done for several initial guesses of the background intensity; and the value that gives the lowest $\\chi^{2}$ between the binned data and the ionization model is chosen as the best fit $J(\\nu_{0})$. However, this is not the optimal technique to use at low redshift where absorption line densities are low. KF93 developed a maximum likelihood technique to address this issue and used it in their measurement of $J(\\nu_{0})$ at $z \\sim 0.5$. However, their measurement was based upon a sample of only 13 QSOs and less than 100 lines, and has correspondingly large error bars. In addition, the value these authors find is lower than the predictions of the models of Haardt \\& Madau (1996), though consistent within the uncertainties, as shown in Figure 13 of Paper II and in Figure~\\ref{fig:lowzcomp} of this paper. Given the importance of the value of the HI ionization rate to the hydrodynamical evolution of the low redshift universe, performing this measurement with a much larger line sample is worthwhile. The low redshift hydrodynamic simulations of Theuns et al.\\ (1998) and Dav\\'{e} et al.\\ (1999) indicate that the evolution of the ionizing background is the primary driver behind the change of character of the Ly-$\\alpha$ forest from high redshift to low redshift, specifically, the break in the number distribution of Ly-$\\alpha$ lines at $z=1.7$ (Morris et al.\\ 1991, Bahcall et al.\\ 1991, Weymann et al.\\ 1998). The growth of structure pulling gas from low density regions into high density regions also contributes to this and other attributes of the evolution of the Ly-$\\alpha$ forest. Shull et al.\\ (1999) estimate the local ionizing background including contributions to the background from starburst galaxies as well as Seyferts and QSOs. Their models include a treatment of the opacity of the low redshift Ly-$\\alpha$ forest using information drawn from recent observational work (Weymann et al.\\ 1998, Penton et al.\\ 2000b). They find that starbursts and AGN could contribute approximately equally to the ionizing background at low redshift, each $\\sim 1.0 \\times 10^{-23}$ ergs s$^{-1}$ cm$^{-2}$ Hz$^{-1}$ sr$^{-1}$. The full HST/FOS archival data set is presented in Paper III and can also be found at \\mbox{\\tt http://lithops.as.arizona.edu/$\\tilde{\\;}$jill/QuasarSpectra} or \\mbox{\\tt http://hea-www.harvard.edu/QEDT/QuasarSpectra}, \\noindent so we describe the data used in this paper only briefly in \\S\\ref{sec-data}. In \\S\\ref{sec-zsys} and \\S\\ref{sec-flux}, we discuss our treatment of two parameters of each sample object which are integral to the proximity effect analysis, systemic redshifts and Lyman limit fluxes. We outline the proximity effect analysis in \\S\\ref{sec-analysis} and we present our results in \\S\\ref{sec-disc}. We discuss the recovery of $J(\\nu_{0})$ from simulated Ly-$\\alpha$ forest spectra in \\S\\ref{sec-sims}. A comparison of the results from radio loud and radio quiet QSOs is given in \\S\\ref{sec-rl}. The maximum likelihood solutions for $J(\\nu_{0})$ found when allowing for an equivalent width threshold that varies across each sample spectrum are discussed in \\S\\ref{sec-varthr}. Solutions for the HI ionization rate are given in \\S\\ref{sec-gam}. The effect of a non-zero $\\Omega_{\\Lambda}$ on our calculations is discussed in \\S\\ref{sec-omegal}, and the effect of the ionizing background on the Ly-$\\alpha$ forest line density is discussed in \\S\\ref{sec-dndz}. We provide comparisons with previous observational work on the low redshift UV background in \\S\\ref{sec-prevres} and with models of this background in \\S\\ref{sec-models}. A discussion of possible systematic effects on this analysis is given in \\S\\ref{sec-systematics}. We conclude with a summary of the results in \\S\\ref{sec-summary}. ", + "conclusions": "\\label{sec-disc} \\subsection{Radio Loudness} \\label{sec-rl} As the results listed in Table~\\ref{table-jnu} indicate, the inclusion of the four blazars and one BL Lac object, all at $z < 1$, in our sample does not change the result significantly. However, there is much observational evidence that radio loud and radio quiet quasars inhabit different environments, namely that radio loud quasars reside in rich clusters while radio quiet quasars exist in galaxy environments consistent with the field (Stockton 1982, Yee \\& Green 1984, 1987, Yee 1987, Yates, Miller, \\& Peacock 1989, Ellingson, Yee, \\& Green 1991, Yee \\& Ellingson 1993, Wold et al.\\ 2000, Smith, Boyle, \\& Maddox 2000). If there is a corresponding increase in the number of Ly-$\\alpha$ absorption lines in the spectra of radio loud objects, this could cause the proximity effect to be suppressed, and the measured log[$J(\\nu_{0})$] to be artificially large. We have therefore divided our sample into radio loud and radio quiet subsamples using the ratio of radio to UV flux to characterize the radio loudness, \\begin{equation} {\\rm RL = log[S(5 \\; GHz)]/log[S(1450} \\; \\mbox{\\AA})]. \\end{equation} The value of RL for each object in our sample is listed in Table~\\ref{tab-flux}. A histogram of these values and the distribution of RL with $z$ for the sample objects are shown in Figure~\\ref{fig:rl}. The division between radio loud and radio quiet was chosen to be RL=1.0. The resulting values of log[$J(\\nu_{0})$] for these subsamples are listed in Table~\\ref{table-jnu}. There is no significant trend for log[$J(\\nu_{0})$] to appear larger for radio loud objects than for radio quiet objects. \\subsection{Non-Zero $\\Omega_{\\Lambda}$} \\label{sec-omegal} We performed the maximum likelihood calculation for the case of a non-zero cosmological constant. This means that the observer-QSO and absorber-QSO luminosity distances that appear in the relationship between $\\omega$ and $z$ (BDO) must be calculated numerically from the expression: \\begin{equation} d_{L} = (1+z) \\frac{c}{H_{0}} \\int_{0}^{z} \\frac{dz^{\\prime}}{E(z^{\\prime})}, \\end{equation} where \\begin{equation} E(z) \\equiv \\sqrt{\\Omega_{\\rm M}(1+z)^{3} + \\Omega_{\\rm k}(1+z)^{2} + \\Omega_{\\Lambda}}, \\end{equation} (Peebles, 1993) as this integral cannot be reduced to an analytical form for $\\Omega_{\\Lambda} \\neq 0$. The calculations in the sections above assume ($\\Omega_{\\rm M}$,$\\Omega_{\\Lambda}$)~=~(1.0,0.0). Here, we perform the maximum likelihood search for $J(\\nu_{0}$) using ($\\Omega_{\\rm M}$,$\\Omega_{\\Lambda}$)~=~(0.3,0.7). For a QSO at $z=0.5$ with a Lyman limit flux density of 0.1 $\\mu$Jy, an absorber at $z=0.48$, and an assumed background of log[$J(\\nu_{0})$]$=-22$., this ($\\Omega_{\\rm M}$,$\\Omega_{\\Lambda}$) results in a value of $\\omega$ that is $\\sim 25$\\% smaller than that inferred in the $\\Omega_{\\Lambda}=0$ case. Unlike all the other solutions performed, we ignore redshift path associated with metal lines and use all redshifts between $z_{\\rm min}^{q}$ and $z_{\\rm max}^{q}$. This does not change the results significantly, but cuts down the computation time substantially. The results are listed in Table~\\ref{table-jnu} and are plotted in Figure~\\ref{fig:lowzcomp}. For comparison, we also give the solutions for $J(\\nu_{0})$ found using the standard parameters, ($\\Omega_{\\rm M}$,$\\Omega_{\\Lambda}$)~=~(1.0,0.0), with this redshift path neglected. We find that ($\\Omega_{\\rm M}$,$\\Omega_{\\Lambda}$)~=~(0.3,0.7), does not change the value of $J(\\nu_{0})$ derived significantly from the value found using ($\\Omega_{\\rm M}$,$\\Omega_{\\Lambda}$)~=~(1.0,0.0). We performed a slightly modified re-analysis of the Scott et al.\\ (2000b) sample of objects at $z \\sim 2$ and found little effect at high redshift as well. The solution found for ($\\Omega_{\\rm M}$,$\\Omega_{\\Lambda}$)~=~(1.0,0.0) was log[$J(\\nu_{0})$]$=-21.09^{+0.20}_{-0.17}$, while for ($\\Omega_{\\rm M}$,$\\Omega_{\\Lambda}$)~=~(0.3,0.7), we find log[$J(\\nu_{0})$]$=-21.25^{+0.20}_{-0.17}$ for these data. \\subsection{$d{\\cal N}/dz$} \\label{sec-dndz} In the case of a size distribution of Ly-$\\alpha$ absorbers that is constant in redshift, the evolution of the number of Ly-$\\alpha$ absorption lines per unit redshift is given by: \\begin{equation} d{\\cal N}/dz={\\cal N}_{0} (1+z)^{2} [\\Omega_{M} (1+z)^{3} + (1-\\Omega_{M}-\\Omega_{\\Lambda}) (1+z)^{2} + \\Omega_{\\Lambda}]^{-0.5}, \\label{equ:noevol} \\end{equation} (Sargent et al.\\ 1980) where ${\\cal N}_{0}$ equals the absorber cross section times the absorber comoving number density times the Hubble distance, $\\pi r_{0}^{2} \\phi_{0} c H_{0}^{-1}$. A plot of $d{\\cal N}/dz$ versus $z$ for non-evolving Ly-$\\alpha$ absorbers in ($\\Omega_{M}$,$\\Omega_{\\Lambda}$)~=~(1.0,0.0) and (0.3,0.7) cosmologies is shown in Figure~\\ref{fig:noevol}. It is clear that non-evolving models are too shallow to fit points at $z > 1.7$, so the normalization is found from a fit to the FOS data. The FOS data at $z < 1.7$ are consistent with a non-evolving population for $(\\Omega_{M},\\Omega_{\\Lambda})=(1.0,0.0)$. The data are less consistent with a non-evolving concordance model in which $(\\Omega_{M},\\Omega_{\\Lambda})=(0.3,0.7)$, though not significantly so. The number density evolution of Ly-$\\alpha$ absorbers over the redshift range $z=0-5$ cannot be approximated with a single power law. There is a significant break in the slope of the line number density with respect to redshift, near $z=1.7$ (Weymann et al.\\ 1998, Paper IV) though Kim, Cristiani, \\& D'Odorico (2001) argue that the break occurs at $z=1.2$. Dav\\'{e} et al.\\ (1999) show from hydrodynamical simulations of the low redshift Ly-$\\alpha$ forest, that the evolution of the line density is sensitive mainly to the HI photoionization rate, but also to the evolution of structure (cf.\\ their Figure 7). The flattening of $d{\\cal N}/dz$ observed by Weymann et al.\\ (1998) is mostly attributed to a dramatic decline in $\\Gamma(z)$ with decreasing $z$. Dav\\'{e} et al.\\ (1999) derive an expression for the density of Ly-$\\alpha$ forest lines per unit redshift as a function of the HI photoionization rate: \\begin{equation} \\frac{d{\\cal N}}{dz} = C [ (1+z)^{5} \\Gamma^{-1}(z)]^{\\beta-1} H^{-1}(z), \\label{equ:dave} \\end{equation} where $C$ is the normalization at some fiducial redshift which we choose to be $z=0$ and $\\Gamma (z)$ can be expressed by Equ.~\\ref{equ:hmgam}. We fit the FOS and MMT absorption line data, binned in $d{\\cal N}/dz$ as presented in Paper IV and Scott et al.\\ (2000a, Paper I), to this function in order to derive the parameters describing $\\Gamma (z)$ implied by the evolution in Ly-$\\alpha$ forest line density. We observe flattening of $d{\\cal N}/dz$ at $z < 1.7$, but not to the degree seen by Weymann et al.\\ (1998) in the Key Project data. As described in Paper IV, we find $\\gamma=0.54\\pm0.21$, for lines above a 0.24~\\AA\\ threshold, while Weymann et al.\\ (1998) measure $\\gamma=0.15\\pm0.23$. See Paper IV for more discussion of the significance and underlying causes of this difference. We find $(A_{\\rm HM},B_{\\rm HM},z_{c},S)=(3.0\\times 10^{-12},0.61,5.5 \\times 10^{-7}, 7.07)$ and $(1.9 \\times 10^{-11}, 0.38, 3.4 \\times 10^{-7}, 6.21)$ for $(\\Omega_{\\rm M},\\Omega_{\\Lambda})=(1.,0.)$ and lines with rest equivalent widths above 0.24 and 0.32~\\AA\\, respectively. These fits to Equ.~\\ref{equ:dave} are shown in Figure~\\ref{fig:dndz}(a). In panel (b), we plot $\\Gamma (z)$, as expressed in Equ.~\\ref{equ:hmgam}, evaluated using the parameters found from the fit to Equ.~\\ref{equ:dave} above. The HM96 solution and the solution derived from the full FOS and MMT data sets are represented by the thick and thin solid lines respectively. The small values of $z_{c}$ derived from $d{\\cal N}/dz$ above translate into ionization rates that do not decrease dramatically with decreasing redshift and result from the less pronounced flattening of $d{\\cal N}/dz$ relative to the Key Project. These fits are particularly insensitive to the normalization, $A_{\\rm HM}$, so the errors on this parameter are large. These fits should therefore not be interpreted as measurements of $\\Gamma (z)$ as reliable as those found directly from the absorption line data. But we find them instructive nonetheless. The observed $\\Gamma (z)$ falls short of the ionization rate needed to fully account for the change in the Ly-$\\alpha$ line density with redshift, indicating that if the value of $\\gamma$ at low redshift is indeed slightly larger than that found by the Key Project, $d{\\cal N}/dz$ may still be consistent with a non-evolving population of Ly-$\\alpha$ absorbers in the sense noted above, but the formation of structure in the low redshift universe must play a significant role in determining the character of the Ly-$\\alpha$ forest line density. \\subsection{Comparison with Previous Results} \\label{sec-prevres} \\subsubsection{Proximity Effect} KF93 performed a similar measurement with a small subsample of this total sample- the HST Quasar Absorption Line Key Project data of Bahcall et al.\\ (1993). We compare our result to that from Sample 2 of KF93, which was constructed from the Bahcall et al.\\ (1993) data excluding one BAL quasar and all heavy element absorption systems. The Key Project sample has since been supplemented (Bahcall et al.\\ 1996, Jannuzi et al. 1998) and those data have been included when appropriate in the complete archival sample of FOS spectra presented in Paper III. The mean intensity KF93 derive from their Sample 2 ($b=35$ km~s$^{-1}$, $\\beta$=1.48, $\\gamma$=0.21) is $5.0^{+20.}_{-3.4} \\times 10^{-24}$ ergs s$^{-1}$ cm$^{-2}$ Hz$^{-1}$ sr$^{-1}$. This result is lower than ours for $z < 1$ by a factor of $\\sim 13$, though the errors are large on both results are large enough that they are consistent. We use 162 lines in our low redshift solution for $J(\\nu_{0})$, 65 more than KF93. \\subsubsection{Direct Measurements} \\label{sec-direct} Several authors have examined the sharp cutoffs observed in the HI disks of galaxies in the context of using these signatures to infer the local ionizing background (Maloney 1993, Corbelli \\& Salpeter 1993, Dove \\& Shull 1994). The truncations are modeled as arising primarily from photoionization of the disk gas by the local extragalactic background radiation field. Using 21 cm observations (Corbelli, Scheider, \\& Salpeter 1989, van Gorkom 1993) to constrain these models, limits on the local ionizing background are placed at $10^{4} < \\Phi_{{\\rm ion}} < 5 \\times 10^{4}$ cm$^{-2}$ s$^{-1}$, where \\begin{equation} \\Phi_{{\\rm ion}} = 2\\pi \\int_{0}^{1} \\mu d\\mu \\int_{\\nu_{0}}^{\\infty} \\frac{J_{\\nu}}{h\\nu} d\\nu = \\frac{\\pi J(\\nu_{0})}{h \\alpha_{s}}, \\end{equation} and where $J_{\\nu}=I_{\\nu}$ for an isotropic radiation field. Additionally, narrow-band and Fabry-Perot observations of H$\\alpha$ emission from intergalactic clouds (Stocke et al.\\ 1991, Bland-Hawthorn et al.\\ 1994, Vogel et al.\\ 1995, Donahue, Aldering, \\& Stocke 1995) place limits of $\\Phi_{{\\rm ion}} \\lesssim 10^{4}$ cm$^{-2}$ s$^{-1}$, or $J(\\nu_{0}) < 7.6 \\times 10^{-23}$ ergs s$^{-1}$ cm$^{-2}$ Hz$^{-1}$ sr$^{-1}$ for $\\alpha_{s}=1.8$, while results from measurements of Galactic high velocity clouds (Kutyrev \\& Reynolds 1989, Songaila, Bryant, \\& Cowie 1989, Tufte, Reynolds, \\& Haffner 1998) imply $\\Phi_{{\\rm ion}} \\lesssim 6 \\times 10^{4}$ cm$^{-2}$ s$^{-1}$, though the ionization of high velocity clouds may be contaminated by a Galactic stellar contribution. Tumlinson et al.\\ (1999) have reanalyzed the 3C273/NGC3067 field using the H$\\alpha$ imaging data from Stocke et al.\\ (1991) as well as new GHRS spectra of 3C273, in order to model the ionization balance in the absorbing gas in the halo of NGC3067. From this analysis, they derive the limits, $2600 < \\Phi_{{\\rm ion}} < 10^{4}$ cm$^{-2}$ s$^{-1}$, or $10^{-23} < J(\\nu_{0}) < 3.8 \\times 10^{-23}$ ergs s$^{-1}$ cm$^{-2}$ Hz$^{-1}$ sr$^{-1}$ at $z = 0.0047$. Weymann et al.\\ (2001) have recently reported an upper limit of $\\Phi_{{\\rm ion}} < 1.01 \\times 10^{4}$ cm$^{-2}$ s$^{-1}$, or $J(\\nu_{0}) < 3.84 \\times 10^{-23}$ ergs s$^{-1}$ cm$^{-2}$ Hz$^{-1}$ sr$^{-1}$ from Fabry-Perot observations of the intergalactic HI cloud, 1225+01, for a face-on disk geometry. If an inclined disk geometry is assumed, this lower limit becomes $J(\\nu_{0}) < 9.6 \\times 10^{-24}$ ergs s$^{-1}$ cm$^{-2}$ Hz$^{-1}$ sr$^{-1}$. These results are summarized in Figure~\\ref{fig:allzcomp}. It is encouraging that the proximity effect value is consistent with the limits on the background set by these more direct estimates which are possible locally. \\subsection{Comparison with Models} \\label{sec-models} Haardt \\& Madau (1996) calculated the spectrum of the UV background as a function of frequency and redshift using a model based on the integrated emission from QSOs alone. The QSO luminosity function is drawn from Pei (1995). The opacity of the intergalactic medium is computed from the observed redshift and column density distributions of Ly-$\\alpha$ absorbers given by Equ.~\\ref{eq:dndzdnh1}. The effects of attenuation and reemission of radiation by hydrogen and helium in Ly-$\\alpha$ absorbers are included in these models. Their result for $q_{0}=0.5$ and $\\alpha_{s}=1.8$ at $z=0$ is $J(\\nu_{0}) = 1.6 \\times 10^{-23}$ ergs s$^{-1}$ cm$^{-2}$ Hz$^{-1}$ sr$^{-1}$. Fardal et al.\\ (1998) compute opacity models for the intergalactic medium (IGM) based on high resolution observations of the high redshift Ly-$\\alpha$ forest from several authors. Shull et al.\\ (1999) extend the models of Fardal et al.\\ (1998) to $z=0$, treating opacity of low redshift Ly-$\\alpha$ forest from observations made with HST/GHRS (Penton et al.\\ 2000a,b) and with HST/FOS (Weymann et al.\\ 1998). Like Haardt \\& Madau (1996), they also incorporate the observed redshift distribution of Lyman limit systems with log(N$_{{\\rm HI}}$) $ > 17$ (Stengler-Larrea et al.\\ 1995, Storrie-Lombardi et al.\\ 1994). Their models also allow for a contribution from star formation in galaxies in addition to AGN. The QSO luminosity function again is taken to follow the form given by Pei (1995) with upper/lower cutoffs at 0.01/10 L$_{*}$. QSO UV spectral indicies are assumed to equal 0.86, while the ionizing spectrum at $\\nu > \\nu_{0}$ has $\\alpha_{s}=1.8$. The contribution to the background from stars was normalized to the H$\\alpha$ luminosity function observed by Gallego et al.\\ (1995) and the escape fraction of photons of all energies from galaxies was taken to be $ =0.05$. The full radiative transfer model described in Fardal et al.\\ (1998) was used to calculate the contribution to the mean intensity by AGN, but not the contribution from stars, as they were assumed to contribute no flux above 4 Ryd, the energies at which the effects of IGM reprocessing become important. These authors find $J(\\nu_{0})=2.4 \\times 10^{-23}$ ergs s$^{-1}$ cm$^{-2}$ Hz$^{-1}$ sr$^{-1}$ at $z\\sim 0$, with approximately equal contributions from AGN and stars, a value somewhat lower than our result for $z < 1$, but which is allowed within the errors. We estimate the contribution to the UV background from star-forming galaxies using the galaxy luminosity function of the Canada-France Redshift Survey (Lilly et al.\\ 1995). At $z \\sim 0.5$, we derive $J^{\\rm gal}(\\nu_{0})= 1.5 \\times 10^{-22}$ ergs s$^{-1}$ cm$^{-2}$ Hz$^{-1}$ sr$^{-1}$, assuming $=1$. The HM96 models for the QSO contribution give $J^{\\rm QSO}(\\nu_{0}) = 5.2 \\times 10^{-23}$ ergs s$^{-1}$ cm$^{-2}$ Hz$^{-1}$ sr$^{-1}$ at $z \\sim 0.5$. These estimates, and the range of measured $J(\\nu_{0})$ in this paper, $\\sim 5-16 \\times 10^{-23}$ ergs s$^{-1}$ cm$^{-2}$ Hz$^{-1}$ sr$^{-1}$ imply an escape fraction of UV photons from galaxies between 4\\% and 70\\%. The $J(\\nu_{0})$ inferred from $d{\\cal N}/dz$ in Section~\\ref{sec-dndz} implies escape fractions well over 100\\%. Bianchi et al.\\ (2001) make updated estimates of the mean intensity of the background with contributions from both QSOs and star-forming galaxies. Their models incorporate various values of the escape fraction of Lyman continuum photons from galaxies which are constant with redshift and wavelength. Our new results at $z < 1.7$ are most consistent with their models of the QSO contribution alone, though some contribution from galaxies, ie. a small f$_{\\rm esc}$, is allowed within the uncertainties. At $z \\sim 3.5$, recent results from Steidel, Pettini, \\& Adelberger (2001) on the Lyman-continuum radiation from high redshift galaxies suggest that these sources become a more important component of the UV background at high redshift. \\subsection{Systematics} \\label{sec-systematics} Drawing on lessons learned from our work on high redshift objects in Paper II, we have made corrections for quasar systemic redshifts before performing the proximity effect analysis, as discussed in \\S\\ref{sec-zsys}. This correction, $\\sim 300$ km~s$^{-1}$, was made to QSO redshifts measured from Ly-$\\alpha$ emission for objects for which no systemic redshift measurement was available. For the low redshifts considered in this paper, redshifts measured from [OIII], MgII, or Balmer emission lines were deemed suitable as QSO systemic redshift measurements. We have removed known gravitational lenses from the sample. As discussed above, we perform the proximity effect analysis omitting and including spectra that show associated absorption and damped Ly-$\\alpha$ absorption and determined that neither of these populations significantly biases our results. Because we are working with low redshift data where line densities are low, we expect that blending has not contributed as strong a systematic effect as in the high redshift sample of Paper II. The curve-of-growth effects discussed in Paper II may still be present, since many lines in the sample have equivalent widths which place them on the flat part of the curve of growth. However, the effects of clustering may be even more important at low redshift than at high redshift. Loeb \\& Eisenstein (1995) showed how the fact that quasars reside in the dark matter potentials of galaxies and small groups of galaxies can influence the proximity effect signature. The peculiar velocities of matter clustered in these potentials can result in Ly-$\\alpha$ absorption at redshifts greater than the quasar emission redshift. We found that including associated absorbers in our sample did not significantly change our results. Recently, Pascarelle et al.\\ (2001) report evidence for a lower incidence of Ly-$\\alpha$ absorption lines arising in the gaseous halos of galaxies in the vicinities of QSOs than in regions far from QSOs. They argue that galaxy-QSO clustering may lead proximity effect measurements to overestimate $J(\\nu_{0})$ at $z < 1$ by a up to a factor of 20. While we agree that most systematic effects in this type of analysis, including clustering, will lead to overestimates of $J(\\nu_{0})$, the agreement between our results and the direct measurements discussed in Section~\\ref{sec-direct} give us confidence that our results are not biased by this large a factor. The hydrodynamic simulations of the low redshift Ly-$\\alpha$ forest of Dav\\'{e} et al.\\ (1999) indicate that, at low redshift, structures of the same column density correspond to larger overdensities and more advanced dynamical states than at high redshift. For a $(\\Omega_{\\rm M},\\Omega_{\\Lambda})~=~(0.4,0.6)$ cosmology, an equivalent width limit of 0.32~\\AA\\ corresponds to an overdensity of $\\sim 1.4$ at $z \\sim 3$, while at $z \\sim 0.6$, this limit corresponds to $\\rho_{H} / \\overline{\\rho_{H}} \\sim 13$. This may have implications on the clustering of Ly-$\\alpha$ absorption lines around QSOs and hence on the values of $J(\\nu_{0})$ derived from the proximity effect. It is possible that we are seeing this clustering effect in the variable threshold solution at $z > 1$, in which the two highest $\\omega(z)$ lines in the sample are responsible for the inability to isolate a reasonable maximum likelihood $J(\\nu_{0})$. We have analyzed a set of 151 QSOs and 906 Ly-$\\alpha$ absorption lines, the subset of the total data set presented in Paper III that is appropriate for the proximity effect. The primary results of this paper are as follows: (1) At low redshift, Balmer, [OIII], and Mg II emission lines are reasonable indicators of QSO systemic redshifts. Ly-$\\alpha$ emission is blueshifted by $\\sim 300$ km~s$^{-1}$ with respect to [OIII]. (2) The value of $J(\\nu_{0})$ is observed to increase with redshift over the redshift range of the sample data, $0.03 < z < 1.67$. Dividing the sample at $z = 1$, we find $J(\\nu_{0})= 6.5^{+38.}_{-1.6} \\times 10^{-23}$ ergs s$^{-1}$ cm$^{-2}$ Hz$^{-1}$ sr$^{-1}$, at low redshift and $J(\\nu_{0})= 1.0^{+3.8}_{-0.2} \\times 10^{-22}$ ergs s$^{-1}$ cm$^{-2}$ Hz$^{-1}$ sr$^{-1}$ at high redshift. (3) The inclusion of blazars at $z < 1$ has no significant effect on the result. There is no significant difference between the values of $J(\\nu_{0})$ derived from radio loud (RL $>$ 1.0) and radio quiet (RL $<$ 1.0) objects, indicating that the observed richness of quasar environments does not distinctly bias the proximity effect analysis. (4) Using information measured and gathered from the literature on each QSO's UV spectral index and solving for the HI ionization rate, yields $1.9 \\times 10^{-13}$ s$^{-1}$ for $z < 1$ and $1.3 \\times 10^{-12}$ s$^{-1}$ for and $z > 1$. Solving directly for the parameters $(A_{\\rm HM},B_{\\rm HM},z_{c},S)$ in the HM96 parametrization of $\\Gamma(z)$ using the HST/FOS data presented by Bechtold et al.\\ (2001) combined with the high redshift, ground-based data presented by Scott et al.\\ (2000a,b) results in $(A_{\\rm HM},B_{\\rm HM},z_{c},S)$~=~($7.6\\times 10^{-13}$, 0.35, 2.07, 1.77) for $\\beta=1.46$ and $(A_{\\rm HM},B_{\\rm HM},z_{c},S)=(3.2\\times 10^{-13}$, 1.45, 2.13, 1.42) for $\\beta=1.7$ for $1.7 < z < 3.8$. (5) Allowing for a varying equivalent width threshold across each QSO spectrum results in consistently higher values of $J(\\nu_{0})$ than are found from the constant threshold treatments. At $z > 1$, the variable threshold solution is not well-constrained. Jackknife experiments indicate that this is due the objects 0743-6719 and 0302-2223, namely the highest $\\omega(z)$ absorption lines in each of their spectra. (6) Allowing for a cosmology in which $(\\Omega_{\\rm M},\\Omega_{\\Lambda}) = (0.3,0.7)$, rather than (1.,0.) has no significant effect on the value of $J(\\nu_{0})$ derived from these data. (7) The $z < 1$ result is in agreement with the range of values of the mean intensity of the hydrogen-ionizing background allowed by a variety of local estimates, including H$\\alpha$ imaging and modeling of galaxy HI disk truncations. To within the uncertainty in the measurement, this result agrees with the one previous proximity effect measurement of the low redshift UV background (KF93). These results are consistent with calculated models based upon the integrated emission from QSOs alone (HM96) and with models which include both QSOs and starburst galaxies (Shull et al.\\ 1999). The uncertainties do not make a distinction between these two models possible. (8) The results presented here tentatively confirm the IGM evolution scenario provided by large scale hydrodynamic simulations (Dav\\'{e} et al.\\ 1999). This scenario, which is successful in describing many observed properties of the low redshift IGM, is dependent upon an evolving $J(\\nu_{0})$ which decreases from $z = 2$ to $z = 0$. However, the low redshift UV background required to match the observations of the evolution of the Ly-$\\alpha$ forest line density is larger than found from the data, indicating that structure formation is playing a role in this evolution as well. Our results and the work of others are summarized in Figure~\\ref{fig:allzcomp}. We find some evidence of evolution in $J(\\nu_{0})$, though it appears that even larger data sets, especially at $z < 1$ and/or improved proximity effect ionization models will be required to improve the significance." + }, + "0112/astro-ph0112223_arXiv.txt": { + "abstract": "HI 21 cm observations with the NRAO 140 Foot telescope have revealed a giant HI cloud (G28.17+0.05) in the Galactic plane that has unusual properties. The cloud is 150 pc in diameter, is at a distance of 5 kpc, and contains as much as 10$^5 {\\rm M_{\\sun}}$ of atomic hydrogen. The cloud consists of a cold core, $T \\sim 10~K$, and a hotter outer envelope, $T \\ge 200~K$. There is no observable difference in the HI line widths, $\\sim 7$ km~s$^{-1}$, between the core and the envelope. Anomalously-excited 1720 MHz OH emission, with a similar line width, is associated with the core of the cloud. The cloud core also exhibits $^{12}$CO and $^{13}$CO self-absorption which indicates that most of the cloud mass is in molecules. The total mass of the cloud is $> 2 \\times 10^5 {\\rm M_{\\sun}}$. The cloud has only a few sites of current star formation. If similar clouds are associated with other observed sites of anomalously-excited 1720 MHz OH emission, there may be as many as 100 more of these objects in the inner galaxy. ", + "introduction": "G28.17+0.05 was discovered in 21cm HI observations made with the NRAO 140 Foot Telescope (Minter et al. 2001). At the time of these observations these data comprised the most detailed look at this part of the Galactic Plane in HI. G28.17+0.05 is shown in Figure 1. It appears as a shell of bright emission surrounding an area of reduced emission centered on ${\\rm l,b} = 28\\fdg17+0\\fdg05$. G28.17+0.05 is estimated to have a distance of $\\sim 5~{\\rm kpc}$ at which its diameter would be $150~{\\rm pc}$. The reduced emission in the central area of G28.17+0.05 (hereafter referred to as the core) is the result of HI self-absorption (see Figure 2). Models of the 21cm HI emission from G28.17+0.05 indicate the temperature in the core is ${\\rm T \\leq 40~K}$ while the shell has ${\\rm T \\geq 200~K}$. The core contains $1/3$ of the total HI mass. The column density of atomic hydrogen through the center is ${\\rm N_{HI} \\geq 5-12 \\times 10^{20}~cm^{-2}}$, depending on the exact value of the temperature in the core. The total atomic hydrogen mass within G28.17+0.0 is $\\sim 10^5 {\\rm M_{\\sun}}$. The line widths (FWHM) of HI in the core and shell are both $\\sim 7~{\\rm km~s^{-1}}$. This results in G28.17+0.05 being gravitationally bound from just its atomic mass alone! \\begin{figure} \\plotfiddle{MinterA2_1.eps}{3in}{0}{50}{50}{-140}{-90} \\caption{G28.17+0.05 21cm HI brightness temperature at $v_{LSR} = 77~{\\rm km~s^{-1}}$.} \\end{figure} \\begin{figure} \\plotfiddle{MinterA2_2.eps}{3in}{0}{50}{50}{-175}{0} \\caption{Spectra of 21cm HI and 1720~MHz OH toward the center of G28.17+0.05. The 1720~MHz OH emission has been scaled by an arbitrary factor so that its structure and position can be compared with the 21cm HI emission. The arrow marks the center of the HI self-absorption.} \\end{figure} ", + "conclusions": "There are several indications that G28.17+0.05 is a young cloud: (1) The emissions lines of HI and OH have a similar width and are dominated by turbulent motions, suggesting that the turbulence in G28.17+0.05 has not yet had time to damp. (2) The lack of star formation. (3) The large amount of atomic matter in G28.17+0.05. We suspect that G28.17+0.05 may the first example of a cloud caught in the transition phase from atomic to molecular as the cloud encounters a spiral arm shock (Minter et al. 2001). Since anomalously excited 1720~MHz OH emission is quite abundant throughout the inner galaxy (Turner 1982), and this part of the galaxy has not been observed in HI with great detail until very recently, there could be many more of these atomic/molecular clouds that have yet to be found or identified. The VGPS (Taylor et al. 2002) will be extremely valuable for studying this object, and may reveal many others like it throughout the Galaxy." + }, + "0112/astro-ph0112015_arXiv.txt": { + "abstract": "In this contribution I am going to present some preliminary results of a high-resolution spectroscopic campaign focussed on the most metal rich red giant stars in $\\omega$~Cen. This study is part of a long term project we started a few years ago, which is aimed at studying the properties of the different stellar populations in $\\omega$~Cen. The final goal of the whole project is the global understanding of both the star formation and the chemical evolution history of this complex stellar system. ", + "introduction": "\\begin{figure} \\plotfiddle{Pancino_fig01.aph.ps}{15cm}{0}{70}{70}{-210}{-70} \\caption{The WFI photometry is displayed. Each panel shows the CMD for one of the CCD chips in the WFI mosaic. The cluster center is roughly located in chip \\#~2. The arrows mark the RGB-a.} \\end{figure} \\begin{figure} \\plotfiddle{Pancino_fig02.aph.ps}{14cm}{0}{70}{70}{-220}{-80} \\caption{{\\it Top Panel:} Zoomed CMD of the upper RGB region. Stars in the RGB-a have been plotted as {\\it small filled triangles}. Large empty symbols show the position in the CMD for stars with known metallicity (from Norris et al. 1996). Different symbols refer to stars with different metallicities: {\\it large open circles} for stars with $-1.5<$[Ca/H]$<-1.3$, {\\it large open stars} for $-1.1<$[Ca/H]$<-0.85$, {\\it large open squares} for $-0.65<$[Ca/H]$<-0.4$, {\\it large open triangles} for [Ca/H]$>-0.3$, respectively. {\\it Bottom Panel:} Histogram of the distribution of the distances from the mean ridge line of the dominant, metal poor RGB. The mean [Ca/H] abundances for the four main components of the RGB are also marked.} \\end{figure} The first surprising result that we got in the study of $\\omega$~Cen came from the wide field photometry that we performed in 1999 at the 2.2 m ESO-MPI telescope (La Silla, Chile), equipped with the WFI (Wide Field Imager) mosaic camera. Figure~1 shows the $(B,B-I)$ CMDs for each of the 8 WFI chips separately. The cluster is roughly centered on chip \\#2. More than 230,000 stars have been plotted: to our knowledge, this is the largest stellar sample ever observed in $\\omega$~Cen. The most striking feature of this CMD is the existence of a {\\it complex} structure in the brighter part of the Red Giant Branch (RGB): at least two main components are visible. Particularly notable is the presence of a narrow sequence (see also Lee et al. 1999), significantly redder and more bent than the bulk of the ``main'' RGB stars, which we call the {\\it anomalous} RGB (hereafter RGB-a). Note that the RGB-a is visible only in the CMD from chip \\#2 (where the cluster center is located) and to a much lesser extent in the nearby chip \\#1 and possibly chip \\#3. The morphology of the RGB-a and its position in the CMD strongly suggest that it is populated by stars much more metal rich than the $\\omega$~Cen bulk population. In order to get a first indication about the metal content of the RGB-a stars, we have combined our catalog with the extensive low resolution spectroscopic survey carried on by Norris et al. (1996), based on the infrared calcium triplet indicator: Figure~2 shows the metallicity distribution of the red giants in $\\omega$~Cen, obtained by using the distance of each star from the mean ridge line of the dominant, metal poor population as a sort of metallicity indicator (Figure~2, Panel {\\it (b)}). We can thus identify three different RGB components, or sub-populations: \\begin{itemize} \\item{{\\it RGB-MP:} the dominant, metal poor population, with [Ca/H]$\\sim-1.4$, which drives the average metallicity of the whole cluster.} \\item{{\\it RGB-MInt:} the intermediate metallicity population, with [Ca/H]$\\sim-1.0$ and a long extended tail reaching up to [Ca/H]$\\sim-0.5$.} \\item{{\\it RGB-a:} the anomalous, newly discovered population, with [Ca/H]$\\sim-0.1$, which represents the extreme metallicity end of the cluster distribution, and comprises $\\sim5\\%$ of the whole RGB population.} \\end{itemize} Thus, we have photometrically isolated a distinct sub-population, the RGB-a, whose sharp and clear-cut shape in our CMD allows us to select a representative sample of stars for high-resolution spectroscopic follow up, in order to obtain their detailed abundance pattern. ", + "conclusions": "The results presented here represent a first step towards the understanding of the star formation history of $\\omega$~Cen. A comprehensive scenario of course requires a much more extensive sample of stars, together with accurate measures of the relative ages and detailed abundance patterns of stars near the Turn-off and Sub Giant Branch regions of the CMD. We are already planning to extend our study in this direction within the framework of the global project presented by Ferraro, Pancino \\& Bellazzini, in this volume." + }, + "0112/astro-ph0112365_arXiv.txt": { + "abstract": "The interstellar medium is turbulent and this induces relative motions of dust grains. We calculate relative velocities of charged grains in a partially ionized magnetized gas. We account for anisotropy of magnetohydrodynamic (MHD) turbulence, grain coupling with magnetic field, and the turbulence cutoff arising from the ambipolar drag. We obtain grain velocities for turbulence with parameters consistent with those in HI and dark clouds. These velocities are smaller than those in earlier papers, where MHD effects were disregarded. Finally, we consider grain velocities arising from photoelectric emission, radiation pressure and H\\( _{2} \\) thrust. These are still lower than relative velocities induced by turbulence. We conclude that turbulence should prevent these mechanisms from segregating grains by size. ", + "introduction": "Dust is an important constituent of the interstellar medium (ISM). It interferes with observations in the optical range, but provides an insight to star-formation activity through far-infrared radiation. It also enables molecular hydrogen formation and traces the magnetic field via emission and extinction polarization. The basic properties of dust (optical, alignment etc.) strongly depend on its size distribution. The latter evolves as the result of grain collisions, whose frequency and consequences depend on grain relative velocities. Various processes can affect the velocities of dust grains. Radiation, ambipolar diffusion, and gravitational sedimentation all can bring about a dispersion in grain velocities. It was speculated in de Oliveira-Costa et al. (2000) that starlight radiation can produce the segregation of different sized grains that is necessary to explain a poor correlation of the microwave and \\( 100\\mu m \\) signals of the foreground emission (Mukherjee et al. 2001). If true it has big implications for the CMB foreground studies. However, the efficiency of this segregation depends on grain random velocities, which we study in this paper. Interstellar gas is turbulent (see Arons \\& Max 1975). Turbulence was invoked by a number of authors (see Kusaka et al. 1970, Volk et al. 1980, Draine 1985, Ossenkopf 1993, Weidenschilling \\& Ruzmaikina 1994) to provide substantial relative motions of dust particles. However, they discussed hydrodynamic turbulence. It is clear that this picture cannot be applicable to the magnetized ISM as the magnetic fields substantially affect fluid dynamics. Moreover dust grains are charged, and their interactions with magnetized turbulence is very different from the hydrodynamic case. This unsatisfactory situation motivates us to revisit the problem and calculate the grain relative motions in magnetized ISM. In what follows, we use the model of MHD turbulence by Goldreich and Sridhar (1995, henceforth GS95), which is supported by recent numerical simulations (Cho \\& Vishniac 2000, Maron \\& Goldreich 2001, Cho, Lazarian \\& Vishniac 2002a, henceforth CLV02). We apply our results to the cold neutral medium (CNM) and a dark cloud to estimate the efficiency of coagulation, shattering and segregation of grains. ", + "conclusions": "\\subsection{Shattering and Coagulation } Consider the cold neutral medium (CNM) with temperature \\( T=100 \\)K, density \\( n_{\\rm H}=30 \\)cm\\( ^{-3} \\), electron density \\( n_{e}=0.045 \\)cm\\( ^{-3} \\), magnetic field \\( B\\sim 1.3\\times 10^{-5} \\)G (Weingartner \\& Draine 2001a, hereafter WD01a). To account for the Coulomb drag, we use the results by WD01a and get the modified drag time \\( t_{drag}=\\alpha t^{0}_{drag} \\). Using the electric potentials in Weingartner \\& Draine (2001b), we get grain charge and \\( \\tau _{L} \\). For the parameters given above, we find that \\( t_{drag} \\) is larger than \\( \\tau _{c} \\) for grains larger than \\( 10^{-6} \\)cm, \\( \\tau _{L} \\) is smaller than \\( \\tau _{c} \\) even for grains as large as \\( 10^{-5} \\)cm. Here, we only consider grains larger than \\( 10^{-6} \\)cm, which carry most grain mass (\\( \\sim 80\\% \\)) in ISM, so we can still use Eq.(\\ref{vpara}) to calculate grain parallel velocities and Eq.(\\ref{vperp}) to get the perpendicular velocity for grain larger than \\( 10^{-5} \\)cm. Nevertheless, the perpendicular velocities of grains smaller than \\( 10^{-5} \\)cm should be estimated as \\( v'_{\\perp }(a)=v_{c}\\times (\\tau _{L}/\\tau _{c})=v_{max}(\\tau _{c}/\\tau _{max})^{1/2}(\\tau _{L}/\\tau _{c})=v_{\\perp }(a)(\\tau _{L}/\\tau _{c})^{1/2}, \\) where \\( v_{\\perp }(a) \\) is given by Eq.(\\ref{vperp}). The results are shown in Fig.1. The critical sticking velocity were calculated in Chokshi et al. (1993)(see also Dominik \\& Tielens 1997).\\footnote{There are obvious misprints in the numerical coefficient of Eq.(7) in Chokshi et al.(1993) and the power index of Young's modulus in Eq.(28) of Dominik \\& Tielens (1997).} However, experimental work by Blum (2000) shows that the critical velocity is an order of magnitude larger than the theoretical calculation. Thus the collisions can result in coagulation for small silicate grains (\\(\\leq 3\\times 10^{-6} \\)cm). With our input parameters, grains do not shatter if the shattering thresholds for silicate is \\( 2.7 \\)km/s as in Jones et al. (1996). Nevertheless, the grain velocities strongly depend on \\( v_{max} \\) at the injection scale. For instance, we will get a cutoff \\( 6\\times 10^{-5} \\)cm due to shattering if \\( v_{max}=10 \\)km/s. For a dark cloud, the situation is different. As the density increases, the drag by gas becomes stronger. Consider a typical dark cloud with temperature \\( T=20 \\)K, density \\( n_{\\rm H}=10^{4} \\)cm\\( ^{-3} \\) (Chokshi et al. 1993) and magnetic field \\( B\\sim 2.3\\times 10^{-4} \\)G. Assuming that dark clouds are shielded from radiation, grains get charged by collisions with electrons: \\( =0.3(r/10^{-5} \\)cm) electrons. The ionization in the cloud is \\( \\chi =n_{e}/n_{tot}\\sim 10^{-6} \\) and the drag by neutral atoms is dominant. From Eq.(\\ref{cutoff}) and the expression for the drag time and the Larmor time, we find \\( \\tau _{L}a \\)) \\begin{equation} \\label{H} \\langle F_{z{\\rm H}}\\rangle =r^{3/2}(r+1)^{1/2}\\gamma (1-y)n_{\\rm H}v_{\\rm H}a^{2}\\left( \\frac{2m_{\\rm H}E}{\\nu }\\right) ^{1/2}, \\end{equation} where \\( r=b/2a, \\), \\( n_{\\rm H}\\equiv n({\\rm H})+2n({\\rm H}_2) \\), \\( y=2n({\\rm H}_2)/n_{\\rm H} \\) is the \\( {\\rm H_2} \\) fraction, \\( \\gamma \\) is the fraction of impinging {\\rm H} atoms and \\( \\nu \\) is the number of active sites over the grain surface. The expected grain velocity is \\( v=\\langle F_{z{\\rm H}}\\rangle t_{drag}/m \\). In the CNM we consider, \\( y=0 \\), adopting the characteristic values in Lazarian \\& Draine (1997), \\( r=1, \\) \\( \\gamma =0.2 \\), \\( E=0.2 \\)eV, and the density of active sites \\( 10^{11} \\)cm\\(^{-2}\\) so that \\( \\nu =80(a/10^{-5} \\)cm\\( )^{2}r(r+1) \\), we get the {}``optimistic{}'' velocity shown in Fig 1. For maximal active site density \\( 10^{15} \\)cm\\( ^{-2} \\), we get the lower boundary of grain velocity \\( v\\simeq 3.3(10^{-5}{\\rm cm} /a)^{1/2} \\)cm/s. The scaling is approximate due to the complexity of coefficient \\( \\alpha \\)(see WD01a Fig.16). Lazarian \\& Draine (1999a,b) have shown that subjected to {\\rm H}\\( _{2} \\) torques alone, grains \\( \\leq 10^{-4} \\)cm should experience frequent thermal flipping, which means that the \\( F_{z{\\rm H}} \\) fluctuates. This flipping results from coupling of grain rotational and vibrational degrees of freedom through internal relaxation and would average out \\( \\langle F_{z{\\rm H}}\\rangle \\). However, the flipping rate depends on the value of the grain angular momentum (Lazarian \\& Draine 1999a). If a grain is already spun up to a sufficient velocity, it gets immune to thermal flipping. Radiative torques (Draine \\& Weingartner 1996) can provide efficient spin if the grain size is comparable to the wavelength. For a typical interstellar diffuse radiation field, the radiative torques are expected to spin up grains with sizes larger than \\( \\sim 4\\times 10^{-6} \\)cm. They will also align grains with rotational axes parallel to the magnetic field. Thus grains should acquire velocities along the magnetic field lines and the corresponding velocities should be compared with those arising from turbulent motions parallel to the magnetic field. It is clear from Fig.1 that for the chosen set of parameters the effect of {\\rm H}\\( _{2} \\) thrust is limited. All in all, we conclude that the radiation effects and H\\(_2\\) thrust are not efficient for segregating grains in typical ISM conditions. \\begin{figure} \\centering \\leavevmode \\resizebox*{0.33\\textwidth}{0.25\\textheight}{\\includegraphics{f1.eps}} \\caption{Grain velocities as a function of radii (solid line) in the CNM. Dashdot line represents parallel velocity due to the drag by compressible modes, dotted line refers to perpendicular velocity from the contribution of the drag by Alfven mode, also plotted is the earlier hydrodynamic result (dashed line). The change of the slope is due to the cutoff of turbulence by ambipolar diffusion. The grain velocity driven by {\\rm H}\\protect\\( _{2}\\protect \\) thrust is plotted to illustrate the issue of grain segregation in the CNM (see text), the part marked by 'o' is nonphysical because thermal flipping is not taken into account. } \\end{figure}" + }, + "0112/astro-ph0112479_arXiv.txt": { + "abstract": "We have searched for nonastrophysical emission lines in the optical spectra of 577 nearby F, G, K, and M main--sequence stars. Emission lines of astrophysical origin would also have been detected, such as from a time--variable chromosphere or infalling comets. We examined $\\sim$20 spectra per star obtained during four years with the Keck/HIRES spectrometer at a resolution of 5 \\kms, with a detection threshold 3\\% of the continuum flux level. We searched each spectrum from 4000~\\AA~-- 5000~\\AA~ for emission lines having widths too narrow to be natural from the host star, as well as for lines broadened by astrophysical mechanisms. We would have detected lasers that emit a power, $P>$50 kW, for a typical beam width of $\\sim$0.01 arcsec (diffraction--limit from a 10-m aperture) if directed toward Earth from the star. No lines consistent with laser emission were found. ", + "introduction": "Most searches for signals from extraterrestrial intelligent civilizations (SETI) have concentrated on radio wavelengths (e.g. Tarter 2001, Cullers 2000, Werthimer et al. 2000, Leigh and Horowitz 2000). However, a preferred wavelength for SETI work has proved elusive. The motivation for this project stemmed from the suggestion of Schwartz and Townes (1961) that advanced civilizations might use optical or near--infrared lasers as an efficient means of interstellar communication (Townes 1998, Lampton 2000). The advantages of optical wavelengths over radio include the higher bandwidth and the narrower diffraction--limited beam size to promote private communication. In addition, the technology needed to produce lasers with detectable powers is not far from current Earth technology (\\S 2.3). Optical SETI (OSETI) projects are currently designed to detect highly energetic, pulsed lasers and continuous--wave lasers. Such efforts include dedicated searches, piggyback searches, and data mining. Many OSETI projects are designed to detect pulsed lasers using fast photometry (Howard \\et~2000, Bhathal 2000, Werthimer \\et~2001). Two or three photodetectors are employed to search for laser pulses having a duration of $\\sim$1 ns in broadband visible light. A high--powered nanosecond pulsed laser will outshine its host star during its brief pulse because at most one photon will be received (for a typical 1-m class telescope) from the star during the 1 ns duration. Alternately, interstellar optical communication may occur via nearly continuous laser emission. In that case, lasers of even moderate power, tens of kilowatts, can outshine the host star when focussed by large apertures during transmission (see \\S 2.3). Here we report a search for ultra narrow emission lines such as predicted by the Schwartz--Townes suggestion. ", + "conclusions": "We would have detected time--varying emission lines from any of the 577 main sequence stars above a threshold of a few percent of the continuum flux. Emission lines of constant intensity would also have been detected had they Doppler--shifted relative to the star during the $\\sim$4 yr of observations. Lines having intrinsic widths from arbitrarily narrow to tens of \\kms would have been detected. No narrow lines consistent with laser emission were detected. However, naturally occurring emission lines were found in the spectra of two stars. Emission lines in the spectra of Solar--type stars are rare between 4000~\\AA~-- 5000~\\AA. Such lines can arise occasionally in the chromospheres or coronae of the most magnetically active stars. In some ($\\sim$20\\%) F, G, K--type main sequence stars, the cores of the Balmer absorption lines are ``filled in'' by a few percent due to chromospheric emission (Herbig 1985, Robinson, Cram, \\& Giampapa 1990). We would have detected temporal variation in that chromospheric filling, if it exceeded a few percent of the continuum level. No such variations were detected in any of our F, G, K--type main sequence stars. This nondetection indicates that such chromospheric variation in the cores of the strong Balmer lines, notably H$\\beta$ and H$\\gamma$, did not exceed a few percent of the continuum level in these stars. Flare stars, most of spectral type M0 or later, more commonly exhibit emission of H$\\beta$. Indeed, we detected time--varying emission of H$\\beta$ in two M--type flare stars. Time--varying emission lines from stars may also arise from infalling volatile material and comets, as is observed for $\\beta$ Pic (see for example Beust et al. 1998). Our search would have detected such astrophysical emission lines if they appeared or varied by more than a few percent of the continuum flux of the star. All such astrophysical origins for emissions lines would yield line widths greater than our resolution of 5 \\kms due to the usual thermal and turbulent broadening, as well as possible rotational and collisional broadening. Hence such lines would be resolved. Nonetheless, no emission lines due to these effects were found in the 577 stars at our threshold of a few percent of continuum flux. The entrance slit of the HIRES spectrometer is rectangular with dimensions 0.57 $\\times$ 3.5 arcsec. At the typical distance of our stars of $\\sim$50 pc, this slit corresponds to 30--175 AU at the star. Thus, sources of emission lines residing within 30 AU of our target stars would have been included in our spectra. Of course, the emission beam must be pointed toward the Keck telescope for detection. Extinction from interstellar dust is negligible toward these target stars due to their proximity ($d<$100 pc). The motivation for this search stemmed from the suggestion of Schwartz and Townes (1961) that advanced civilizations might communicate from one colony to another within a planetary system, or from one planetary system to another, by using optical lasers as the carrier. The advantages of optical communication over radio include its higher bandwidth and its smaller diffraction beam to promote private communication. A nondetection of extraterrestrial intelligence carries no value unless it excludes a plausible model of life in the universe. One might imagine the following model. The Milky Way Galaxy has an age of $\\sim$10 Gyr, while only 4.6 Gyr were required to spawn a technological species on Earth. Approximately half of the stars in the Galactic disk are older than the Sun, and they formed with comparable amounts of heavy elements. Approximately 50\\% of these old disk stars do not have a stellar companion within 100 AU (Duquennoy and Mayor 1991), thereby permitting stable planetary orbits. Therefore, $\\sim$25\\% of the stars in the Galactic disk have requisite characteristics of adequate age, chemical composition, and dynamical quiescence to serve as sites for the development of technological species. The anthropocentric nature of these stellar criteria are apparent, and thus serve merely as a useful guide for a model. Coincidentally, the target stars used here are systematically middle--aged or older and void of close companions. They were selected for the Doppler planet search at Keck from among stars in the Solar neighborhood. They were chosen to have ages greater than 2 Gyr, as judged from the CaII H\\&K chromospheric emission line, because of the resulting photospheric stability of older stars. Moreover the target stars have no known stellar companions within 2 arcsec (typically $\\sim$50 AU) to prevent any companion's light from entering the spectrometer slit (but see Vogt et al. 2002 and Nidever et al. 2002 for a few recently detected, but negligibly bright, close companions). Moreover, spectroscopic analyses of the stellar spectra show that 90\\% of the target stars contain amounts of heavy elements that are within a factor of 2 of the Solar value ([Fe/H] = -0.3 to +0.3; D.A.Fischer, private communication). Thus, the majority of the 577 target stars for this SETI search are older than the Sun, have roughly Solar chemical composition, and have no perturbing stellar companions within 50 AU. This stellar sample is clearly biased. Thus our observations test a limited model of the Galaxy in which older, single stars harbor technological civilizations. The speculative model posits that some fraction of these target stars contain civilizations that sent probes to, or established colonies on, planets, moons or other platforms either in their host planetary system or in other systems. This model of behavior, albeit anthropocentric, stems from the tendencies of Homo Sapiens toward exploration. To communicate with such probes or colonies over distances of Astronomical Units to parsecs, some fraction of these civilizations might be modeled as using electromagnetic waves, including optical lasers. We model the solid angle subtended by laser beams as follows. Some fraction, $f_{\\rm tech}$, of the total number of target stars, $N_{\\rm stars}$=577, harbor technological civilizations. On average each such civilization emits some number of laser beams, N$_{\\rm beams}$, perhaps in arbitrary directions. We consider that each beam subtends an average solid angle $\\Omega_{\\rm beam}$. The total solid angle, $\\Omega_{\\rm Total}$, subtended by all arbitrarily oriented beams is \\begin{equation} \\Omega_{\\rm Total} = N_{\\rm stars} f_{\\rm tech} N_{\\rm beams} \\Omega_{\\rm beam} \\end{equation} \\noindent We ignore here details of overlap of the beams, their actual distribution of power and beam size, and any purposeful aiming toward or away from us. Any properties of the lasers that are constructed with the Earth as a consideration could drastically change the probability of interception of such beams. In our survey, we have monitored $N_{\\rm stars}$= 577 stars. The nondetection obtained in the survey here suggests that the total solid angle subtended by lasers emitting 50 kW or more is constrained as $\\Omega_{\\rm Total} < 4\\pi$ sr . Thus for arbitrarily directed beams, we constrain the model as \\begin{equation} f_{\\rm tech} N_{\\rm beams} \\Omega_{\\rm beam} < 4\\pi / 577 \\end{equation} The nondetection reported here suggests lasers emitting detectable power must satisfy the above constraint for the combination of $f_{\\rm tech}$, $N_{\\rm lasers}$, $\\Omega_{\\rm beam}$. Apparently, advanced civilizations communicating by kilowatt optical lasers are not so common as to fill 4$\\pi$ sr. As a touchstone, we consider the case of the diffraction--limited laser with an aperture of 10-m (Keck--to--Keck), giving a solid angle for one beam, $\\Omega_{\\rm beam} \\approx 10^{-15}$ sr. For transmitters of such aperture, a laser emitting 50 kW would be detectable. For this case, each star would have to emit $\\sim10^{13}$ arbitrarily oriented beams on average, each with power $P \\ge$ 50 kW, in order that one would likely be oriented, serendipitously, toward us. Clearly, the observed nondetection implies that such numerous, narrow beams, however implausible, do not exist in sufficient numbers to permit likely detection. These observations offer a poor constraint on lasers of arbitrary orientation, as their solid angles are simply too small in this model. One might consider a Galactic model consisting of wider laser beams, enhancing the chances of interception. In that case, the present nondetection would impose a constraint of fewer lasers emanating from the target stars, but would require proportionally higher power per laser for detection at 6--$\\sigma$ above the stellar continuum. More meaningfully, the nondetection reported here suggests that none of the 577 target stars harbors a civilization that has purposefully directed a laser with $P \\ge $50 kW toward Earth (for the case of a 10--meter diffraction--limited transmitting aperture). This modest detection threshold of 50 kW highlights the ease with which an advanced civilization could signal us, if desired. But such is apparently not the case." + }, + "0112/astro-ph0112153_arXiv.txt": { + "abstract": "We present the first detection of the mean flux of the optical extragalactic background light (EBL) at 3000, 5500, and 8000\\AA. Diffuse foreground flux at these wavelengths comes from terrestrial airglow, dust--scattered sunlight (zodiacal light), and dust--scattered Galactic starlight (diffuse Galactic light). We have avoided the brightest of these, terrestrial airglow, by measuring the absolute surface brightness of the night sky from above the Earth's atmosphere using the Wide Field Planetary Camera2 (WFPC2) and Faint Object Spectrograph (FOS), both aboard the Hubble Space Telescope (HST). On the ground, we have used the duPont 2.5\\,m Telescope at Las Campanas Observatory (LCO) to obtain contemporaneous spectrophotometry of ``blank'' sky in the HST field of view to measure and then subtract foreground zodiacal light from the HST observations. We have minimized the diffuse Galactic light in advance by selecting the HST target field along a line of sight with low Galactic dust column density, and then estimated the low--level Galactic foreground using a simple scattering model and the observed correlation between thermal, 100\\micron\\ emission and optical scattered flux from the same dust. In this paper, we describe the coordinated LCO/HST program and the HST observations and data reduction, and present the resulting measurements of the EBL. Galaxies brighter than $V=23$\\,\\ABmag\\ are not well sampled in an image the size of the WFPC2 field of view. We have therefore measured the EBL from unresolved and resolved galaxies fainter than $V=23$\\,\\ABmag\\ by masking out brighter galaxies. We write as EBL23 to emphasize this bright magnitude cut--off. From absolute surface photometry using WFPC2 and ground--based spectroscopy, we find mean values for the EBL23 of 4.0 ($\\pm2.5$), 2.7 ($\\pm1.4$), and 2.2 ($\\pm1.0$) in units of \\tto{-9}\\escsa\\ in the F300W, F555W, and F814W bandpasses, respectively. The errors quoted are $1\\sigma$ combined statistical and systematic uncertainties. The total flux measured in resolved galaxies with $V>23$\\,\\ABmag\\ by standard photometric methods is roughly 15\\% of the EBL23 flux in each band. We have also developed a new method of source photometry, uniquely suited to these data, with which we can measure the ensemble flux from detectable sources much more accurately than is possible with standard methods for faint galaxy photometry. Using this method, we have quantified systematic biases affecting standard galaxy photometry, which prevent light from being recovered in isophotes within a few percent of the sky level. These biases have a significant effect on faint galaxy counts. The flux from resolved sources as measured by our ensemble photometry method is 3.2 ($\\pm$ 0.22), 0.89 ($\\pm$ 0.01), and 0.76 ($\\pm$ 0.01) in units of \\tto{-9} \\escsa\\ in the F300W, F555W, and F814W bandpasses, respectively, with $1\\sigma$ combined errors. These values, the total flux from resolved sources, represent {\\it absolute minima} for the EBL23 in each band, and are roughly 30\\% of the mean flux we measure for the total EBL23. ", + "introduction": "The Extragalactic Background Light (EBL) is the integrated, mean surface brightness of both resolved and unresolved extragalactic sources. In wavelength range of these observations (2500--9500\\AA), the EBL is dominated by flux from stellar nucleosynthesis at redshifts $z\\lta 9$, and also includes flux from gravitational potential energy such as accreting black holes (Fabian \\& Iwasawa 1999), brown dwarfs, gravitationally collapsing systems, and possibly decaying particles. As such, the EBL is the fossil record of the star formation history of the universe and a fundamental measure of the luminous energy content of the universe. As we show in Figure \\ref{fig:foregrndflux}, upper limits from previous attempts to measure the EBL and lower limits from integrated galaxy counts constrain the EBL to an expected level of roughly $1\\times10^{-9}$ \\escsa\\ at 5500\\AA, or 28.2 AB mag arcsec$^{-2}$.\\footnote{All surface brightnesses are specified in \\escsa\\ unless specifically noted otherwise. AB mag is defined as AB mag$=-2.5\\log F_\\nu - 48.6$, as usual, with F$_\\nu$ given in \\esch.} Because the combined flux from foreground airglow, zodiacal light, and diffuse galactic light (also plotted in Figure \\ref{fig:foregrndflux}) is at least 100 times brighter than the EBL, a detection of the EBL requires measurement accuracies of better than 1\\%. Previous attempts to measure the optical EBL have employed a variety of different approaches. Mattila (1976) attempted to isolate the EBL by differencing integrations on and off the line of sight to ``dark'' clouds at high Galactic latitude, under the assumption that the clouds acted as a ``blank screens,'' spatially isolating all foreground contributions from the background. This pioneering work produced upper limits and identified both the rapid temporal variability of terrestrial airglow and extinction and the spatial variability of diffuse Galactic light, which proved to be the primary obstacles to these early efforts to measure the EBL. Toller (1983) later attempted to avoid both atmospheric and zodiacal foregrounds by using data taken with the {\\it Pioneer 10} spacecraft at a distance of 3 AU from the Sun, beyond the zodiacal dust cloud. Poor spatial resolution ($2^{\\circ}$), however, prevented the accurate subtraction of discrete Galactic stars, let alone the diffuse Galactic component from these data. Dube, Wilkes, \\& Wilkinson (1977, 1979) made the first effort to measure and subtract foreground contributions explicitly based on geometrical modeling of airglow and Galactic foregrounds but including spectroscopic measurement of the zodiacal light (ZL) flux by a technique similar to that which we have adopted. Rapid variability caused uncertainty in their airglow subtraction which dominated the errors in their results. \\begin{figure}[t] \\begin{center} \\includegraphics[width=3.5in]{bfm1_fig1.ps} \\caption{\\footnotesize Relative surface brightnesses of foreground sources, upper limits on the EBL23 (see \\S\\ref{intro}), and lower limits based on the integrated flux from resolved galaxies ($V_{555}>23$ AB mag) in the HDF (Williams \\etal 1996). The spectral shape and mean flux of zodiacal light and of diffuse galactic light (DGL) are shown at the levels we detect in this work. The airglow spectrum is taken from Broadfoot \\& Kendall (1968) and is scaled to the flux level we observe at 3800--5100\\AA\\ (see \\S\\ref{zl}). The effective bandpasses for our HST observations are indicated at the bottom of the plot.} \\label{fig:foregrndflux} \\end{center} \\end{figure} In this work, we take advantage of the significant gains in technology and in understanding of the foreground sources which have been achieved since the last attempts to measure the optical EBL (see Mattila 1990 for a review). The most significant technological advance is panoramic, linear CCD detectors. Those on-board HST allowed us to completely avoid bright, time--variable airglow and provide sub--arcsecond spatial resolution. High spatial resolution allowed us to resolve stars to $V\\sim 27.5$ mag and thereby eliminate the possibility of significant contamination from unidentified Galactic stars in the field. Ground--based spectrophotometry with CCDs also made possible much more accurate measurement of the foreground zodiacal light than could be achieved with narrow--band filters and photometers, as were used by Dube, Wickes, \\& Wilkinson (1977, 1979). Finally, IRAS has provided maps of the thermal emission from dust at high Galactic latitudes. We have use the IRAS maps to select a line of sight for these observations which has a low column density of Galactic dust in order to minimize the DGL contribution caused by dust--scattered starlight, and also to estimate the low--level DGL which cannot be avoided. Our measurement of the EBL utilizes three independent data sets. Two of these are from HST: (1) images from the Wide Field Planetary Camera 2 (WFPC2) using the F300W, F555W, and F814W filters, each roughly 1000\\AA\\ wide with central wavelengths of 3000, 5500, and 8000\\AA, respectively; and (2) low--resolution spectra (300\\AA\\ per resolution element) from the Faint Object Spectrograph (FOS) covering 3900--7000\\AA. The FOS data were taken in parallel observing mode with the WFPC2 observations. While flux calibration of WFPC2 images and FOS spectra achieve roughly the same accuracy for point source observations, the increase in spatial resolution, $10^4$ times larger field of view, lower instrumental background, and absolute surface brightness calibration achievable with WFPC2 make it better suited than FOS to an absolute surface brightness measurement of the EBL. Nonetheless, the FOS observations do provide a second, independent measurement of the total background flux of the night sky, also free of terrestrial airglow and extinction, but with greater spectral resolution than the WFPC2 images. The third data set consists of long--slit spectrophotometry of a region of ``blank'' sky within the WFPC2 field of view. These data were obtained at the 2.5m duPont telescope at Las Campanas Observatory (LCO) using the Boller \\& Chivens spectrograph simultaneously with one visit of the HST observations (6 of 18 orbits). The flow chart in Figure \\ref{fig:flowchart} summarizes the reduction procedures for each data set used in this measurement, the results obtained from each data set individually, and the coordination of those results to produce a measurement of the EBL. In this paper, we begin by describing the foreground sources in \\S\\ref{foreg} and the details of HST scheduling in \\S\\ref{sched}. The observations and data reduction of both HST data sets, WFPC2 and FOS, are discussed in detail in this paper. WFPC2 observations and data reduction are discussed in \\S\\ref{wfpc2}. In \\S\\ref{wfpc2.backg}, we present the first results from the WFPC2 data, which are measurements of the total sky flux (foregrounds plus background) in each bandpass. The FOS observations, data reduction, and results are discussed in \\S\\ref{fos} and \\S\\ref{fos.resul}. The modeling of diffuse Galactic light is discussed in \\S\\ref{dgl}. The LCO data and measurement of ZL are discussed in Bernstein, Freedman \\& Madore (2001b, henceforth Paper II). A summary of that work is given in \\S\\ref{zl}. In \\S\\ref{ebl.minim}, we present a measurement of the minimum flux of the EBL from resolved sources, which can be made using the WFPC2 data alone. The implications of that result are also discussed in \\S\\ref{ebl.minim}. Finally, in \\S\\ref{ebl.detec} we combine the results of the individual data sets and modeling (horizontal connections shown in the flow chart) to obtain a measurement of the EBL. The implications of these results are discussed in Bernstein, Freedman \\& Madore (2001c, henceforth Paper III). ", + "conclusions": "} \\tablehead{ \\colhead{Source} &\\colhead{Bandpass}&\\colhead{Data source} & \\colhead{Flux} & \\colhead{Random} &\\colhead{Systematic }} \\startdata Total & F300W \t& {\\sc hst/wfpc2} & 33.5\t\t&$(\\pm$ 4.9\\%)\t&$[\\pm$ 5.6\\%]\\nl \\ \\ Background & F555W \t& {\\sc hst/wfpc2} & 105.7 \t&$(\\pm$ 0.3\\%) \t&$[\\pm$ 1.4\\%]\\nl & F814W \t& {\\sc hst/wfpc2} & 72.4 \t\t&$(\\pm$ 0.2\\%) \t&$[\\pm$ 1.4\\%]\\nl & F555W\\tablenotemark{a} & {\\sc hst/fos} & 111.5\t\t&$(\\pm$ 0.7\\%) \t&$[\\pm$ 2.8\\%]\\nl Zodiacal & 4600--4700\\AA& {\\sc lco} & 109.4\t\t&$(\\pm$ 0.6\\%)\t&[$\\pm$ 1.1\\%]\t\\nl \\ \\ Light & F300W \t& {\\sc lco}\\tablenotemark{b} & 28.5\t\t&$(\\pm$ 0.6\\%)\t&[-1.1\\%,+1.2\\%] \\nl & F555W \t& {\\sc lco}\\tablenotemark{b} & 102.2 \t&$(\\pm$ 0.6\\%)\t&[-1.1\\%,+1.1\\%] \\nl & F814W \t& {\\sc lco}\\tablenotemark{b} & 69.4\t\t&$(\\pm$ 0.6\\%)\t&[-1.3\\%,+1.1\\%] \\nl Diffuse & F300W & {\\sc dgl} model & 1.0\t& \\nodata \t& [+25\\%,-50\\%] \\nl \\ \\ Galactic & F555W & {\\sc dgl} model & 0.8\t& \\nodata \t& [+25\\%,-50\\%] \\nl \\ \\ Light & F814W & {\\sc dgl} model & 0.8\t& \\nodata \t& [+25\\%,-50\\%] \\nl \\enddata \\tablecomments{All fluxes are in units of 1\\tto{-9}\\escsa. For a source with constant flux in $ F_{\\lambda}$, filters F300W, F555W, and F814W have effective wavelengths $\\lambda_0(\\Delta\\lambda)$ =3000(700), 5500(1200), and 8100(1500)\\AA. For a source with a solar spectrum, effective wavelengths are $\\lambda_0$ = 3200, 5500, 8000\\AA.} \\tablenotetext{a}{Observed FOS spectrum, convolved with the WFPC2/F555W bandpass to allow direct comparison with the WFPC2 results.} \\tablenotetext{b}{LCO measurement of zodiacal light, extrapolated to the WFPC2 bandpass by applying a correction for changing zodiacal light color with wavelength relative to the solar spectrum. The zodiacal light flux through the WFPC2 bandpasses was identified using SYNPHOT models, the uncertainty due to which is included in the uncertainty for the filter calibration and is shared with the systematic uncertainty for the total background flux.} \\end{deluxetable} \\begin{deluxetable}{lcccr} \\tablewidth{30pc} \\tablecaption{EBL Results and Uncertainties \\label{tab:cum.errors}} \\tablehead{ \\colhead{Bandpass}& \\colhead{Random} & \\colhead{Systematic}& \\colhead{Combined}& \\colhead{EBL($\\pm1\\sigma$)}\\\\ \\colhead{} & \\colhead{$\\sigma_R$ (68\\%)} & \\colhead{$\\sigma_S$ (68\\%)} & \\colhead{$\\sigma$ (68\\%)} & \\colhead{} } \\startdata \\multicolumn{5}{l}{\\underline{Detected EBL23 (WFPC2 + LCO)}\\tablenotemark{a}} \\nl F300W\\hspace{0.5in} & 2.1\t& 1.5 & 2.5\t& 4.0 $(\\pm 2.5)$ \\nl F555W & 0.6\t& 1.3 & 1.4\t& 2.7 $(\\pm 1.4)$ \\nl F814W & 0.4\t& 0.9\t\t& 0.0\t& 2.2 $(\\pm 1.0)$ \\nl \\multicolumn{5}{l}{\\underline{Minimum EBL (WFPC2)}\\tablenotemark{a}} \\nl F300W & 0.19\t& 0.13\t\t& 0.22\t& 3.2 $(\\pm 0.22)$ \\nl F555W & 0.003\t& 0.009\t\t& 0.01\t& 0.89 $(\\pm 0.01)$ \\nl F814W & 0.002\t& 0.007\t\t& 0.01\t& 0.76 $(\\pm 0.01)$ \\nl \\multicolumn{5}{l}{\\underline{Detected EBL23 (FOS + LCO)}\\tablenotemark{a}} \\nl F555W & 0.7\t& 2.7 & 2.8\t& 8.5 $(\\pm 5.6)$ \\nl \\multicolumn{5}{l}{\\underline{Flux from detected sources in HDF ($m>23$ AB mag)}} \\nl F300W &\t&\t\t&\t& 0.66\t \t\\nl F450W &\t&\t\t&\t& 0.51\t \t\\nl F606W &\t&\t\t&\t& 0.40\t \t\\nl F814W &\t&\t\t&\t& 0.27\t \t\\nl \\multicolumn{5}{l}{\\underline{Published number counts}\\tablenotemark{b}}\\nl \\multicolumn{3}{l}{F300W ($1827.5$\\,AB\\,mag is 0.57\\tto{-10}\\escsa\\ at $V_{555}$ as measured by standard photometric methods. In keeping with the discussion in the previous section, we estimate that only 35\\% of the light is recovered from galaxies with $27.723$\\,AB\\,mag to be 8.9\\tto{-10}\\escsa. We emphasize that this estimate of the minimum EBL23 is indeed a minimum estimate from which sources will be excluded due to the surface brightness biases and detection limits of our own images as well as the HDF (see Figures \\ref{fig:detec.rabmumumag} and \\ref{fig:detec.hdfmumumag}). Incompleteness due to surface brightness detection limits is discussed further in Paper III. Following the same method for F300W and F814W data leads to the minimum EBL23 values summarized in Table \\ref{tab:cum.errors}. The total combined statistical and systematic error for this minimum EBL23 measurement is roughly $\\pm 1$\\tto{-11} {\\escsa}, $\\times 100$ smaller than the error for the EBL23 detections because the conversion to physical units can occur after foregrounds are subtracted. Comparing the total flux from detected galaxies measured using ensemble photometry versus standard methods, the photometry errors affecting standard methods are clearly worse at UV wavelengths where signal--to--noise ratio is lower, calibration is less accurate, and the instrumental PSF is broader than in the other filters." + }, + "0112/astro-ph0112403_arXiv.txt": { + "abstract": "Some hot, massive, population-I Wolf-Rayet (WR) stars of the carbon subclass are known to be prolific dust-producers. How dust can form in such a hostile environment remains a mystery. Here we report the discovery of a relatively cool, extended, multi-arc dust envelope around the star WR112, most likely formed by wind-wind collision in a long-period binary system. We derive the binary orbital parameters, the dust temperature and the dust mass distributions in the envelope. We find that amorphous carbon is a main constituent of the dust, in agreement with earlier estimates and theoretical predictions. However, the characteristic size of the dust grains is estimated to be $\\sim 1 \\mu m$, significantly larger than theoretical limits. The dust production rate is $6.1\\times10^{-7} M_\\odot {\\rm yr^{-1}}$ and the total detectable dust mass is found to be about $2.8 \\times 10^{-5} M_\\odot$ (for $d=4.15$ kpc). We also show that, despite the hostile environment, at least $\\sim$20\\% of the initially-formed dust may reach the interstellar medium. ", + "introduction": "Classes of objects known to produce significant quantities of dust include asymptotic giant-branch stars, red giants, novae, supernovae (each of which contributes about equal rates, $\\sim 10^{-3} \\, M_\\odot {\\rm yr}^{-1}$, of dust in the whole Galaxy), along with planetary nebulae and proto-stars (each of which yields $\\sim$10 times less; Dwek 1985). Surprisingly, some hot, massive, population I WR stars also form dust \\citep{wdh87}. WR stars are the evolved descendants of massive O-type stars. They consist mainly of He-burning cores surrounded by hot envelopes that drive fast, dense winds with average mass-loss rates $\\sim 10^{-5} \\, M_\\odot {\\rm yr}^{-1}$ and terminal velocities $v_\\infty \\sim 1000 - 4000 \\, km s^{-1}$. There are three successive WR phases, WN, WC and WO, characterized by the dominant emission lines of N, C, and O, respectively, in their optical spectra. All dust-making WR stars belong to the carbon-rich, hydrogen-poor WC subclass \\citep{wil95}. The WR dust-makers are remarkable for two main reasons: (1) the absolute rate of formation is very high, up to $10^{-6} \\, M_\\odot {\\rm yr}^{-1}$ in dust alone, and occurring in either a periodic (due to enhanced wind-wind compression at or near periastron passage in long-period WC + O binaries with eccentirc orbits) or sustained fashion (in single late-type WC stars or moderately short-period WC + O binaries with circular orbits); and (2) the dust is formed in a hot, extremely hostile environment, where the formation process is still unknown. Here we report on near and mid-infrared imaging observations of the dust envelope surrounding the star WR 112 (spectral class WC9). The morphology of this envelope provides clues to the nature of the stellar system, while the photometry allows us to estimate the properties of the dust and the total dust mass. ", + "conclusions": "The recently acquired high-resolution, high signal-to-noise near-mid-IR images of the WR star WR112 have enabled us to resolve and, for the first time, study in detail the extended dust envelope thought to be formed in the wind-wind collision zone of the long-period, $P \\approx 25 \\, {\\rm yr}$, binary. A simple approach involving a minimum of assumptions has allowed us to derive the basic dust properties: We find that the dust emissivity approximately follows the trend expected for amorphous carbon. The radial temperature-distribution shows the dust envelope to be in thermal equilibrium. The dust production-rate corresponds to $\\sim 6\\% $ of the total mass-loss rate, if for the latter we assume a typical value of $10^{-5} \\, M_\\odot \\, {\\rm yr}^{-1}$. The characteristic size of the dust particles, $a \\sim 1\\mu m$, turns out to be much larger than expected from state-of-the-art models of dust growth. We also show that, despite the harsh conditions, $\\sim 20\\%$ of the newly formed dust can survive and escape the system, thus enriching the interstellar medium. This may have a direct implication for the process of heavy-element enrichment in the early Universe, when massive stars completely dominated the stellar scene." + }, + "0112/astro-ph0112290_arXiv.txt": { + "abstract": "The remarkable efficiency of the UVES spectrograph at the VLT has made it possible to push high-resolution, high-S/N ground observations of the \\lya forest down to $z \\sim 1.5$, gaining new insight into the physical conditions of the intergalactic medium and its evolution over more than 90\\% of the cosmic time. The universal expansion, the UV ionizing background and the gravitational condensation of structures are the driving factors shaping the number density and the column density distribution of the absorbers. A (limited) contribution of UV photons produced by galaxies is found to be important to reproduce the observed evolutionary pattern at very high and low redshift. The Lyman forest contains most of the baryons, at least at $z>1.5$, and acts as a reservoir for galaxy formation. The typical Doppler parameter at a fixed column density is measured to slightly increase with decreasing redshift, but the inferred temperature at the mean density is increasing with redshift. The signatures of HeII reionization and feedback from the formation of galactic structures have possibly been detected in the Lyman forest. ", + "introduction": "A sample of 8 QSOs with $1.7 < z_{\\rm em} < 3.7$ has been observed \\cite{kim01,kim02} with VLT/UVES at a typical resolution $45000$ and S/N $\\sim 40-50$. Thanks to the two-arm design of the spectrograph \\cite{UVES}, a high efficiency has been achieved in the whole optical range, from the atmospheric cutoff to $1 \\mu m$, which translates immediately in the possibility of obtaining new results on the Lyman forest, especially at $z \\mincir 2.5$. The data have been reduced with the UVES pipeline \\cite{UVESpip} - an non-negligible factor in maximizing the scientific output per unit time - and analyzed with the VPFIT package \\cite{VPFIT}. ", + "conclusions": "" + }, + "0112/astro-ph0112259_arXiv.txt": { + "abstract": "We report on photometry and high-resolution spectroscopy for IRAS\\,19285+0517. The spectral energy distribution based on visible and near-IR photometry and far-IR fluxes shows that the star is surrounded by dust at a temperature of $T_{\\rm {d}}$ $\\sim$ 250~K. Spectral line analysis shows that the star is a K giant with a projected rotational velocity $v\\,sin i$ = 9 $\\pm$ 2 km s$^{-1}$. We determined the atmospheric parameters: $T_{\\rm {eff}}$ = 4500 K, log $g$ = 2.5, $\\xi_{t}$ = 1.5 km s$^{-1}$, and [Fe/H] = 0.14 dex. The LTE abundance analysis shows that the star is Li-rich (log $\\epsilon$(Li) = 2.5$\\pm$0.15), but with essentially normal C, N, and O, and metal abundances. Spectral synthesis of molecular CN lines yields the carbon isotopic ratio $^{12}$C/$^{13}$C = 9 $\\pm$3, a signature of post-main sequence evolution and dredge-up on the RGB. Analysis of the Li resonance line at 6707~\\AA\\ for different ratios $^{6}$Li/$^{7}$Li shows that the Li profile can be fitted best with a predicted profile for pure $^{7}$Li. Far-IR excess, large Li abundance, and rapid rotation suggest that a planet has been swallowed or, perhaps, that an instability in the RGB outer layers triggered a sudden enrichment of Li and caused mass-loss. ", + "introduction": "Observations of lithium in the atmospheres of red giants continue to pose theoretical challenges. Stars evolving up the red giant branch from the main sequence develop a deep convective envelope that dilutes the surface lithium abundance. Thus, red giants are predicted to have a lithium abundance considerably below the initial or interstellar abundance of log $\\epsilon$(Li)$\\approx$3.0. Classical calculations (Iben 1967) predict a reduction by a factor of 30 and 60 for solar metallicity stars of 1~$M_\\odot$ and 3~$M_\\odot$, respectively. If some lithium has been destroyed prior to evolution off the main sequence, the red giant's lithium abundance will be even further reduced. However, as a very luminous asymptotic red giant (4-7 M$\\odot$), theory and observations show that surface lithium can be replenished and even increased above the initial or interstellar abundance thanks to hot bottom burning in intermediate-mass stars (Sackmann \\& Boothroyd 1992). This paper is concerned with a red giant that has yet to evolve to the asymptotic giant branch but that has a lithium abundance greatly in excess of that expected of a star on the red giant branch. Wallerstein \\& Sneden (1982) discovered the first such red giant: HD\\,112127 with $\\log \\epsilon$(Li) $\\simeq$ 3.0, i.e., the interstellar value. In the last two decades, additional examples have been discovered; Charbonnel \\& Balachandran (2000) list 17 stars (including 3 subgiants, 2 early-AGB stars, and 2 for which evolutionary status has yet to be determined) with $\\log \\epsilon$(Li) $\\geq$ 2.0, and include several additional examples with a lower lithium abundance which is most probably in excess of that expected of a red giant. Three of the 17 stars have a lithium abundance clearly greater than the stars' probable initial value. Charbonnel \\& Balachandran (2000) lists $v\\, sin i$ values for 22 giants of which 10 have $v\\, sin i$ $>$ 8 km s$^{-1}$. If we include the PDS~365 (Drake et al. 2001) and the IRAS~19285+0517 (present study), the number of Li-rich giants for which $v\\, sin i$ $>$ 8 km s$^{-1}$ rises to 12. It now appears that more than half of the Li-rich giants are rotating at a unusually high rate for normal K giants for which average $v\\,sin i \\sim$ 2.0 (de Medeiros et al. 2000). This proportion drops to a mere 2\\% for the more common slowly rotating K giants (Drake et al. 2001). Most of them exhibit a pronounced infrared excess (Gregorio-Hetem et al. 1992; Gregorio-Hetem, Castilho, \\& Barbuy 1993). Here, we report an analysis of a K giant with large infrared excess, first reported to be a Li-rich candidate by de la Reza et al. (1997). The star is IRAS 19285+0517, also known as PDS100 after the Pico dos Dias Survey (de la Reza, Drake, \\& da Silva 1996). No prior quantitative spectroscopy of this star has been reported. We show from high-resolution echelle spectra that the star is Li-rich ($\\log \\epsilon$(Li) $\\simeq$ 2.5) and rapidly rotating ($v\\,sin i$ $\\simeq$ 9 km s$^{-1}$). ", + "conclusions": "With the discovery of extra solar planets in close elliptical orbits, accretion of giant planets by red giants is a virtual certainty. That this phenomenon can account for all of the known Li-rich giants is uncertain. Accretion without activation of the $^{3}$He reservoir certainly cannot account for the super Li-rich stars which have more Li than the ISM value. It would be helpful to extend the detailed analyses of the Li-rich giants. Extensions should include a uniform abundance analysis for the light elements (and isotopes) C, N, and O for a sample including normal K giants, so that subtle differences between Li-rich and Li-normal giants may be distinguished. Extant analyses show that there are no large differences (Berdyugina \\& Savanov 1994; da Silva et al. 1995). Measurement of the beryllium abundance in Li-rich giants would be valuable. Accretion of a planet or brown dwarf restores the Be abundance. Conversion of the $^{3}$He reservoir to $^{7}$Li does not return the Be abundance to its pre-diluted value (Castilho et al. 1999). Measurements of the $^7$Li/$^6$Li ratio should be extended to other stars. Positive detection of $^{6}$Li would favor the accretion scenario. In this regard, we note the recent detection of $^6$Li in the atmosphere of a main sequence star with a giant extra solar planet (Israelian et al. 2001) suggests that the star has accreted one or more former planets. Early accretion of this kind will not account for Li-rich giants like IRAS 19285+0517, because in evolution to the red giant phase, the convective envelope will dilute the lithium to below the observed abundance of IRAS\\,19285+0517." + }, + "0112/astro-ph0112545_arXiv.txt": { + "abstract": "{We present a detailed and systematic investigation of correlated spectral and timing properties of the Z source GX 349+2, using a huge amount $(\\sim 40$~ks) of good quality data from September 29, 1998 to October 13, 1998, obtained with the Proportional-Counter-Array on-board RXTE satellite. We, for the first time, give a detailed comparison between the normal-branch-properties and the flaring-branch-properties of this source, as it shows an extended normal branch (which is a rare phenomenon for it) for our case. In our work, the properties of the peaked noise are analyzed as functions of the position on the Z-track (in color-color-diagram), and they are discussed in connection with the peaked noise seen for another Z source Cyg X-2 (at low overall intensities). This will help to construct theoretical models for the peaked noise, as well as to understand the physics behind the Z-shaped tracks traced by these kind of sources in color-color-diagram. We also find a QPO (centroid frequency $\\sim 3.8$ Hz) at the upper part of the flaring branch. This is for the first time, a QPO is seen from this source and hence is very important in understanding the nature of its X--ray emitting components. ", + "introduction": "The bright low-mass X-ray binary (LMXB) GX 349+2 (also called Sco X-2) belongs to a class, called Z sources (Hasinger and van der Klis 1989). These are the most luminous X--ray binaries, which are believed to contain neutron stars as the accreting objects, since two of them viz, Cyg X-2 (Smale 1998) and GX 17+2 (Kuulkers et al. 2001) have exhibited type I X-ray bursts, which are the characteristic of the neutron star. A Z source traces out a Z-shaped track in X--ray color-color diagram (CD) and hardness-intensity diagram (HID). The Z-track generally consists of three parts, horizontal branch (HB), normal branch (NB) and flaring branch (FB). It is generally believed that the inferred-mass-accretion-rate increases along the Z-track from HB to FB (Hasinger et al. 1990). So far six Z sources have been discovered, which are further divided into two subclasses: (1) Cyg-like: Cyg X-2, GX 5-1 \\& GX 340+0, and (2) Sco-like: Sco X-1, GX 349+2 \\& GX 17+2 (Kuulkers et al. 1994, 1997). Sco-like objects have smaller and slanted HBs and larger FBs, while HBs for Cyg-like sources are comparatively larger and horizontal, and their FBs are generally much smaller. The long term variation in shape and position of Z-track (secular motion) has been observed for Cyg-like source (Kuulkers et al. 1994, 1996; Kuulkers and van der Klis 1996). It has been suggested that Cyg-like sources are being viewed at higher inclination angle compared to Sco-like sources (Kuulkers et al. 1994) and contain neutron stars of higher magnetic field strength (Psaltis et al. 1995). Quasi-periodic-oscillations (QPO) and noise components are found in the power spectra of Z sources. These features are generally well-correlated with the position of the source on the Z-track (Hasinger and van der Klis 1989; van der Klis 1995). There are three types of common noise: very-low-frequency-noise (VLFN), low-frequency-noise (LFN) and high-frequency-noise (HFN). The noise components VLFN and HFN are common in all the three branches, but LFN is observed only in HB. QPOs with frequencies in the range $15-60$~Hz are generally observed in HB and in the upper parts of NB (van der Klis 1995). These are called horizontal-branch-oscillations (HBO). A QPO with the frequency in the range $5-8$~Hz is observed from the middle part of NB up to the NB/FB vertex and is called normal-branch-oscillation (NBO). A sudden increase in centroid frequency of NBO is observed at the NB/FB vertex and according to the common belief, NBO transforms to FBO (flaring-branch-oscillation) at this point (Dieters and van der Klis 2000). As the source moves up along FB, both centroid frequency and full-width-half-maximum (FWHM) of FBO increases. In addition to the low frequency QPOs, kHz QPOs $(200-1200$~Hz) are also observed for all the Z sources (see van der Klis 2000, for a review). GX 349+2 is very similar to Sco X-1 in many respects. For example, both objects exhibit strong flaring behavior and the orbital periods are also similar $(P_{\\rm orb} \\sim 18.9$~hr for Sco X-1 and $\\sim 22$~hr for GX 349+2). However, some properties of GX 349+2 are different from those observed for the other five Z sources. First, it never showed a horizontal branch. Besides, instead of NBO and FBO (together called N/FBO), a broad peaked noise with a centroid frequency and FWHM of around 6 Hz and 10 Hz respectively were observed in its FB (EXOSAT observation; Ponman et al. 1988). It was found that the width of the peaked noise component decreases with increasing intensity, which is contrary to the suggestion made by several models (Lamb et al. 1985; Boyle et al. 1986; Hameury et al. 1985). Ponman et al. (1988) noticed that the strength of the peaked noise is maximum in the intermediate intensity band, i.e., at the NB/FB vertex. However in their work, the lower part of FB and NB could not be differentiated clearly, since they divided the data according to the intensity and not according to the position along the Z curve. They also investigated energy dependence of peaked noise properties and found that rms strength of peaked noise was higher in the higher energy bands and there was no time lag between hard and soft photons. Observation with GINGA Large-Area-Counter indicates that the width and the centroid frequency of peaked noise component does not vary significantly as the source moves along FB (O'Neill et al. 2001). They found that the strength of the peaked noise are maximum in the lower part of FB $(\\sim 10\\%$ of the way up the FB) and it becomes weaker as the source moves up the FB. Kuulkers \\& van der Klis (1998) reported the detection of a similar peaked noise component at the lower part of FB using $\\sim 4$~hrs of RXTE observation. A broader and somewhat weaker peaked noise was detected when the source was in the NB. They also showed that rms strength of peaked noise increases with increasing photon energy. Till now for GX 349+2, a narrow NBO or FBO has never been observed. Inspite of being unique (Kuulkers \\& van der Klis 1998) among the Z sources, GX 349+2 is the least-understood one, and even not well-observed. In this paper, for the first time we present a detailed quantitative study of this source with good quality RXTE PCA data. A similar peaked noise (as observed in GX 349+2) was observed for the Cyg-like Z source Cyg X-2 at low overall intensities (Kuulkers et al. 1999). We therefore compare the Z curves for these two sources, in order to understand the physics behind the peaked noise. In the absence of a proper model, our detailed study is expected to help formulate one. In section 2, we describe the observations and the procedure of analysis. We summarize the results in section 3 and discuss the implications in section 4. ", + "conclusions": "In this paper, we have analyzed $\\sim 14$~days' (total good-time duration is $39.8$~ks) data for the Z source GX 349+2 observed by RXTE PCA. We have computed color-color diagram (CD) and hardness-intensity diagram (HID) for this source with $256$~s time average. Such a big time average has been taken in order to differentiate between the normal branch (NB) and the flaring branch (FB) in an unambiguous way, and even with it, we get enough number of points in CD and HID, as we have a huge good-time data set. We have got an extended normal branch for this source, which for almost all of the earlier observations, showed a very small normal branch (like a blob). Therefore, this is for the first time, it has been possible to compare the properties of NB and FB (and also to study the importance of NB/FB vertex) for this source with long-term good quality data. To understand the nature of the property-variation of the source more quantitatively, we have defined rank number $(S_z)$ along the Z-track (as described in section 2). We have chosen $S_z = 2$ for the NB/FB vertex and $S_z = 3$ for the end point of FB. This is to facilitate the comparison with other five Z sources (that show horizontal branch), as for these sources HB/NB vertex is generally chosen as $S_z = 1$ and NB/FB vertex as $S_z = 2$. Our analysis shows that for $\\sim 79\\%$ of total observation time, the source remained in the FB (see Figure 3 and Table 1). This is in accordance with the fact that GX 349+2 was almost always found to be in FB. The Z-track given in Figure 1 is actually a combination of several Z-tracks, traced by the source in different days. Figure 4 shows how the source moves in CD in the time scale of $\\sim 1$ hour on a particular day. Here it moves from FB to NB (covers $\\sim 17$\\% of the length of the Z-track in $\\sim 30$ minutes), i.e., in the direction of decreasing `inferred-accretion-rate'. We also see (from Table 1) that the source can move in both the directions along the Z-track. We have found two kinds of noises in the power spectra: (1) very-low-frequency-noise (VLFN) and (2) peaked noise (PN). For all the $S_z$ values (i.e., whole of the Z-track), VLFN has been found (like another Sco-like source GX 17+2; Homan et al. 2001). We have fitted the VLFN by a power law $(0.1-1.0$~Hz). The index-value comes in the range $1.3-2.4$ for NB and $1.1-2.8$ for FB (these values were $0.4-1.5$ and $0.8-4.0$ for GX 17+2 respectively; Homan et al. 2001). Therefore, we see that the VLFN is much steeper in NB for GX 349+2. The corresponding percentage rms for our source varies in the range $2\\%-13.1\\%$ in NB and $0.8\\%-5.9\\%$ in FB (according to Homan et al. 2001, for GX 17+2 these values were $0.4\\%-0.8\\%$ and $0.5\\%-1.6\\%$ respectively, for the same range of frequency). Therefore, in general VLFN for GX 349+2 is much stronger. The earlier studies have shown that VLFN strength for this source increases along FB (Ponman et al. 1988; Paul et al. 2001), but for our case it does not vary systematically along this branch (see Table 2). The broad QPO (that is to some extent similar to N/FBO of other Z sources) was found in FB of GX 349+2 for all the past observations. By the definition of QPO (quality factor $Q > 2)$, this feature in power spectra may not be called a QPO. O'Neill et al. (2001) called it `FBN', as they found it mostly in the flaring branch (they had only one $S_z$-point in NB). However we call it `PN', as we have discovered it in both NB and FB, and its nature does not change much from one branch to another. We, for the first time, have studied the properties of PN quantitatively (i.e., using the $S_z$ parameter) for both NB and FB. O'Neill et al. (2001) attempted it with GINGA data, but their data-quality was poor and they did not get an extended normal branch (as is mentioned earlier). As a result, they gave the ranges of centroid frequency $(\\nu_c$; $4-7$ Hz), full-width-half-maximum (FWHM; $6-12$ Hz) and the maximum value of percentage rms $(\\sim 3\\%)$ of PN for the full Z-track. Kuulkers \\& van der Klis (1998) did not study the variation of PN-properties with $S_z$. Therefore, they have found the values of $\\nu_c$, FWHM and percentage rms in NB $(\\sim 9.4$~Hz, $\\sim 16$~Hz and $3\\%$ respectively) and in FB $(\\sim 6$~Hz, $\\sim 11$~Hz and $4\\%$ respectively), and not the ranges. In Table 2, we have displayed a detailed variation of these parameters with $S_z$. For our case, $\\nu_c$ of PN comes in the range $5.0-8.3$~Hz (on average $6.5$~Hz) for NB and $4.7-6.6$~Hz (on average $5.5$~Hz) for FB. Therefore $\\nu_c$ on average slightly decreases, as the source moves from NB to FB. This is in accordance with the earlier results (Kuulkers \\& van der Klis 1998). The ranges of FWHM come out to be $5.5-11.2$~Hz (for NB) and $4.8-11.0$~Hz (for FB), which is also not very different from what previous studies found. From Table 2, we see that the ranges of percentage rms are $2.3\\%-5.3\\%$ in NB and $1.6\\%-4.4\\%$ in FB. Therefore, percentage rms in general slightly decreases from NB to FB, which is contrary to the results of O'Neill et al. (2001) and Kuulkers \\& van der Klis (1998). However, roughly in accordance with the results of O'Neill et al. (2001), we have found that the peak value of percentage rms (in FB) occurs at $\\sim 15\\%$ of the way up the FB. Table 2 shows that PN appears upto 80\\% of the length of FB, whereas previous results (O'Neill et al. 2001) show that it occurs upto atmost 40\\% of FB-length. It is also possible that the variation of PN-properties depends on the direction of motion of the source along the Z-track. For example, for observation no. 3 (source moves from NB to FB) FWHM increases from 5.5 Hz to 8 Hz, and for observation no. 13 (source moves in opposite direction, i.e., from FB to NB) FWHM increases from 9 Hz to 11 Hz. Therefore, although these two observations were made on different days, it is likely that the variation of PN-properties along the Z-track is not reversible. It is worth comparing the results of GX 349+2 with those of other two Sco-like sources, namely, GX 17+2 and Sco X-1. We have found that Sco X-1 has the highest value of $l_{\\rm NB}/l_{\\rm FB}$ (ratio of NB-length to FB-length) among these three sources and GX 17+2 has the lowest value (see Homan et al. 2001; Dieters \\& van der Klis 2000). GX 349+2 have never traced HB in the CD unlike the other two sources. But the most important difference is (as mentioned in section 1) that GX 17+2 and Sco X-1 show narrow QPOs (NBO in the normal branch and FBO in the flaring branch), whereas GX 349+2 shows a broad QPO-like structure (PN in both NB and FB). There is a possibility that these two phenomena have similar origin. This is because they have a few properties common. For example, $\\nu_c$ for NBO has the value $(5-7$~Hz) similar to that for PN, and both N/FBO and PN can be fitted by Lorentzian. Besides, the strength of each of N/FBO (Sco X-1; Dieters \\& van der Klis 2000) and PN (GX 349+2; see Table 3) increases with increasing photon energy. In addition to that, there is no time-lag between high-energy and low-energy photons for both N/FBO and PN. However, there are several mismatch among the properties of these two. It is already mentioned that N/FBO is narrow $(Q$-value almost always is greater than 2), while PN is broad $(Q$-value is generally less than 1). In addition to that, NBO and FBO are two different QPOs, although it is widely believed that these two are same phenomena and NBO quickly transforms to FBO at the NB/FB vertex. However, there is no doubt that PN in both NB and FB is the same phenomenon. Even if NBO and FBO are the same QPO (i.e., N/FBO), there are convincing differences between it and PN. At the NB/FB vertex, NBO-frequency suddenly increases by a factor of $\\sim 2$ to become FBO-frequency, and the values of $\\nu_c$ and FWHM of FBO then increase along the flaring branch, while the values of these parameters for PN remain almost same throughout the Z-track. Therefore, the properties of the source that cause the track in CD to turn at the NB/FB vertex, is not probably connected to PN, whereas they must be correlated to N/FBO. Furthermore, for GX 17+2, NBO occurs upto 35\\% of the length of NB from the NB/FB vertex and FBO appears upto 20\\% of the length of FB from the same vertex (Homan et al. 2001), and these numbers are 50\\% and 10\\% respectively (Dieters \\& van der Klis 2000) for Sco X-2. But for GX 349+2, we see that PN is present in the whole length of NB and in the 80\\% length of FB. It is therefore very likely that N/FBO and PN have different origins. However, for one value (2.64) of $S_z$, we have found (see Figure 9) a narrow PN (or QPO with $Q > 2)$. This may indicate some connection between N/FBO and PN, if this QPO is actually a PN (see a latter paragraph for discussion). From the content of the previous paragraph, it may be clear that a single model for both N/FBO and PN is very difficult to formulate. According to the standard model for N/FBO (Fortner et al. 1989; Lamb 1991; Psaltis et al. 1995), radial oscillations in the optical depth of radial inflow (caused by the radiation pressure at near Eddington luminosity) produce a rocking in the X--ray spectrum, which gives rise to NBO. Although this model may be able to explain N/FBO of Sco X-1, it can not explain PN, as PN-frequency does not increase with the `inferred accretion rate' in FB. Besides for our case, PN has been observed 80\\% of the way up the FB. According to the standard model, oscillations are supposed to be suppressed at such high accretion rate. An alternative model (Titarchuk et al. 2001), that identifies NBO-frequency as the spherical-shell-viscous-frequency, is also not adequate to explain PN. A peaked noise (similar to that for GX 349+2) was observed for the Cyg-like Z source Cyg X-2 (Kuulkers et al. 1999) at low overall intensities. The noise-component extends from 2 Hz to 20 Hz in the power spectrum, peaking near $6-7$~Hz (percentage rms $3\\%)$. To investigate whether the peaked noise components of GX 349+2 and Cyg X-2 have the same origin, we have computed the CD of GX 349+2 using the same energy ranges as given in Kuulkers et al. (1999). We have found (Figure 11) that at low overall intensities, the Z-track of Cyg X-2 looks similar to that of GX 349+2 (none of them shows HB), and although for GX 349+2, the Z-track is shifted towards higher soft color, there may be some connection between these two sources through the hard color. This is because the hard-color-value of the NB/FB vertex for Cyg X-2 increases (i.e., shifts towards (actually becomes more) that for GX 349+2), when it changes from high (or, medium) overall intensity to low overall intensity (the gradual change has not been observed). All these indicate that the nature of the X--ray emitting components of GX 349+2 and that of Cyg X-2 (at low overall intensities) may be similar to some extent and the intensity in the energy range $6.4-16.0$~keV (used to calculate hard color) may be an important clue to understand such similarity. If the origin of peaked noise components of the two sources are actually same, then it may be possible that such noise components are originated in the region, which produces most of the luminosity in $6.4-16.0$~keV energy range. This is supported by the fact that the rms strength of PN is higher in the energy range $7-16$~keV compared to that in a lower energy range (see Table 3). However, as the CD represents a rough spectral behavior, to establish a connection between the two PN components, it is essential to compare the energy spectra of these two sources in details. We have already mentioned that for $S_z = 2.64$, we have found a QPO for GX 349+2. This is, to our knowledge, the first low-frequency-QPO seen for this source (if we do not consider PN as QPO). The FBOs with $\\nu_c \\ge 10$~Hz have been observed from other five Z sources. But the QPO we see for this source has a much lower frequency $(\\nu_c = 3.8$~Hz) and therefore, it is certainly not a FBO. There is a possibility that it has the same origin as of PN (because at the time when it occurs, PN does not occur), but appears with higher quality factor. However from Figure 6, we see that its properties (marked by a `square') are very different from the properties of PN. Therefore it is unlikely that these two have the same origin. A similar 26 Hz QPO has been observed from Cyg X-2 during the intensity dip (Kuulkers \\& van der Klis 1995). There, the authors suggest that a thick torus like structure, formed due to puffing up of inner disk at large accretion rate may obscure the inner emission region from our line of sight and oscillation in such a obscuring torus may provide the explanation for this type of the QPOs. This model requires a higher inclination of the source. The inclination of Cyg X-2 is $\\sim 60^\\circ$ (Orosz and Kuulkers 1999). If the QPO seen for GX 349+2 is also due to the oscillation in such a thick torus, then the inclination of this source should be close to that of Cyg X-2. As mentioned in section 1, GX 349+2 is not a well-observed source and certainly is the most poorly understood one among the Z sources. However, it shows interesting phenomenon like PN (which is probably also observed from Cyg X-2 for a particular position of Z-track in CD). Therefore, the study of this source may be very important to understand the physics behind the tracks traced by the Z sources in CD and HID. The detailed study of GX 349+2 in this paper will be useful for this purpose, as well as will help to formulate a correct theoretical model for PN." + }, + "0112/astro-ph0112284_arXiv.txt": { + "abstract": "We performed high--resolution simulations of two stellar collisions relevant for stars in globular clusters. We considered one head-on collision and one off-axis collision between two 0.6 $M_{\\odot}$ main sequence stars. We show that a resolution of about 100 000 particles is sufficient for most studies of the structure and evolution of blue stragglers. We demonstrate conclusively that collision products between main-sequence stars in globular clusters do not have surface convection zones larger than 0.004 $M_{\\odot}$ after the collision, nor do they develop convection zones during the `pre-main-sequence' thermal relaxation phase of their post-collision evolution. Therefore, any mechanism which requires a surface convection zone (i.e. chemical mixing or angular momentum loss via a magnetic wind) cannot operate in these stars. We show that no disk of material surrounding the collision product is produced in off-axis collisions. The lack of both a convection zone and a disk proves a continuing problem for the angular momentum evolution of blue stragglers in globular clusters. ", + "introduction": "It has been shown that the products of main-sequence -- main-sequence collisions appear in the colour-magnitude diagrams of clusters as blue stragglers \\cite{SLBDRS97,SBH97}. Since blue stragglers are readily observable in clusters, they form an ideal population with which to probe the dynamical evolution of the cluster. The dynamical state of a globular cluster (its density profile, velocity dispersion, amount of mass segregation etc.) will determine the rate and nature of the collisions which occur in the cluster. As the cluster evolves, the kinds of collision that occur will change. Therefore, the population of collision products in a cluster can be used to probe the history of the cluster \\cite{SBEG00,HTAP01}. However, in order to use collision products in this way, there are two issues which must first be understood. First, we know that blue stragglers can also be formed through the merger of two components of a binary system. These blue stragglers will probably have different properties than those formed from collisions. We must either be confident that the population we are observing is collisional in origin (e.g. from cluster density considerations), or be able to distinguish between the two populations. Secondly, we also need to be sure that we understand the formation and evolution of the collision products themselves. This paper is concerned with addressing the second point. In this paper, we present the highest resolution smoothed particle hydrodynamic (SPH) simulations of collisions between main-sequence stars to date. Most recent computations have $\\sim 10^4$ particles \\cite{LRS96,SBH97}, with the highest resolution simulation using $10^5$ particles \\cite{SFLRW01}. In this paper, we increase the number of particles to $10^6$. There are three main reasons to extend this kind of simulation to such high resolution. The first, and simplest, is to make sure that no fundamental change in our understanding of blue stragglers occurs. In earlier work, Benz \\& Hills \\shortcite{BH87} performed an SPH simulation using 1024 particles, and concluded that collision products are fully mixed (i.e. the resulting star is chemically homogeneous). Ten years later, these simulations were repeated but with a factor of 10-50 more particles \\cite{LRS96}. Because of the higher resolution, it became clear that collision products are NOT chemically homogeneous, but instead retain some memory of the chemical profile of the parent stars of the collision. We feel confident that such a fundamental change in our understanding will not happen again when we increase the resolution, but we want to be sure. More importantly, we performed these high resolution calculations to answer two fundamental questions about the structure of the collision products immediately after the collision. The first involves the structure of the outer layers of the star, and the second is concerned with the material which is thrown off by the star during the collision itself. There has been some debate in the literature recently about whether the collision product has a surface convection zone shortly after the end of the collision. The presence of such a convection zone could have ramifications for both the surface chemical abundances and the rotation rate of the star when it reaches its main sequence \\cite{LL95}. Previous SPH simulations and subsequent evolution calculations have shown that no convection zone exists at the end of the collision, nor does it appear during the initial thermal relaxation of the star \\cite{LRS96,SLBDRS97}. However, members of the community argue that the previous simulations are unable to resolve the outermost regions of the star well at all, due to their low particle number and use of equal-mass particles (Livio, private communication). In this paper, we resolve this issue by increasing the resolution of our simulations in the outer regions of the parent stars in particular, and by using variable mass particles. We are also interested in following the material which is thrown off from the stars during the collision. We have discovered, using previous generation simulations, that blue stragglers which are formed by an off-axis collision have an angular momentum problem \\cite{SFLRW01}. Namely, these stars retain too much of their angular momentum, and have no apparent way of losing it during their thermal relaxation phase after the collision. As a result, the collision products spin up during their collapse to the main sequence, and inevitably rotate faster than their break-up velocity. Such collision products can never become blue stragglers, since they tear themselves apart before they reach the main sequence. We wish to follow the lost material with greater resolution, to study the amount of angular momentum that this material carries away with it. We also wish to follow the outermost material of the parent stars to see if some of it forms a disk around the collision product. If so, then we can plausibly suggest that magnetic locking to a disk is responsible for removing angular momentum from the collision product. In \\S\\ref{method} we describe the methods used to model the stellar collisions and their subsequent evolution. We present our results in \\S\\ref{results}, and discuss their implications in \\S\\ref{summary}. ", + "conclusions": "\\label{summary} We have performed the highest resolution SPH simulations of collisions between main-sequence stars to date, using up to $10^6$ SPH particles. The collisions were between equal mass main-sequence stars in globular clusters -- the stars have masses of 0.6 $M_{\\odot}$, metallicity Z=0.001, age of 15 Gyr and a relative velocity at infinity of 10 km/s. We varied the resolution of our simulations from $10^4$ to $10^6$ SPH particles, and we looked at a head-on collision and an off-axis collision. These simulations took between $\\sim 6$ CPU hours and $\\sim 5000 $ CPU hours, and the CPU time scaled approximately linearly with N $\\log$ N. We had three goals for performing this research. Our first goal was to determine the optimal resolution for SPH simulations of this kind. Since we want to use the results of the SPH simulations to study the detailed structure and evolution of the collision products, we need to have an accurate model of the distribution of the SPH particles and their properties. On the other hand, we do not want to waste valuable computer resources going to a higher resolution than is necessary. We have shown that both head-on and off-axis simulations with $\\sim$ 100 000 particles give essentially the same information as runs with higher resolution, and therefore we suggest that most simulations with similar requirements can confidently use on order $10^5$ particles. Our second goal was to use these high resolution simulations to study the outer few percent of the collision products. In particular, we were interested in the existence of a surface convection zone, for two reasons. One consequence of a convection zone is mixing of elements in the convection zone. If the zone reaches deep enough in a star to dredge up nuclear processed material, for example, the surface abundances of elements like C, N and O will be different from what we would expect for primordial material. In blue stragglers and other collision products, surface abundances could be used to determine, perhaps, the masses of the parent stars. However, this can only be done if we understand whether the abundances should be mixed by convection or not. The second consequence of a surface convection zone is angular momentum loss. Stars can lose angular momentum through a magnetic wind \\cite{K88}. This process is most effective in stars with deep convection zones, since they can support stronger magnetic fields and since the angular momentum is drained out of the entire (uniformly rotating) convection zone. We have shown that the product of the head-on stellar collision investigated in this paper does not have a surface convection zone larger than 0.004 $M_{\\odot}$. For comparison, the Sun is considered to have a fairly shallow convection zone, with a mass of 0.02 $M_{\\odot}$ \\cite{GDPK92}. A convection zone less than about 0.01 $M_{\\odot}$ is certainly not deep enough to substantially modify the surface abundances of any element, with the possible exception of lithium, beryllium and boron (which are destroyed at low temperatures). Angular momentum loss is also going to be very small for stars with such small convection zones. Our final goal was to determine whether the off-axis stellar collision products had a circumstellar disk. If a rotating star has a disk with a mass of about 0.01 $M_{\\odot}$, and it is locked to the disk by a magnetic field, then angular momentum will be transported from the star to the disk by the field \\cite{K91}. The net effect of the angular momentum transport will be to slow the star's rotation rate. This effect is thought to operate in pre-main-sequence stars \\cite[and references therein]{SPT00}. Unfortunately, our simulations do not show any evidence for a disk. Based on previous work \\cite{BH87}, we expect that blue stragglers which were formed from collisions with larger impact parameters will have circumstellar disks. However, the low impact parameter collision we modelled in this simulation is still rotating too quickly to contract to the main sequence without losing some angular momentum. Sills et al. \\shortcite{SFLRW01} showed that blue stragglers that are formed by stellar collisions have an angular momentum problem. The collision products are formed with large total angular momentum, and some of them are rotating near their break-up velocity. The collision products are large objects, and they start to contract to the main sequence. As they contract, they spin up since they have no way to lose any angular momentum. They begin to reach break-up velocities and become unstable. Even by shedding material, they cannot remove enough angular momentum to remain bound, and they tear themselves apart. We are forced to one of two conclusions; either blue stragglers are not made by stellar collisions (since head-on collisions are very rare, and all off-axis collisions have too much angular momentum), or collision products have some way of losing their angular momentum. The two standard ways stars are thought to lose angular momentum is through a magnetic wind, or through disk-locking. We have shown that both these scenarios are not viable, since these collision products do not have surface convection zones or circumstellar disks immediately after the collision. However, it is interesting to note the work of Durisen et al. \\shortcite{DGTB86}, who show that if a polytrope is rotating very rapidly, some outer material is thrown off into a disk, leaving a stable (but non--axisymmetric) central object. In order for this to happen, the star has to have a ratio of rotational kinetic energy to gravitational energy $\\beta \\geq 0.3$. At the end of the collision, our off-axis collision product has $\\beta \\sim 0.045$, and is therefore stable against any bar instability. However, as the star evolves and contracts, it is possible that $\\beta$ could becomes high enough to trigger some of these hydrodynamic rotational instabilities. In this paper, we have limited ourselves to two very specific collisions. The two stars involved in the collision were of equal mass, and we only investigated two choices of impact parameter. There have been other studies of main-sequence star stellar collisions which covered more parameter space \\cite{BH87,LRS96,SBH97}. These studies have clearly and carefully outlined the properties of collisions under different circumstances. These high resolution simulations were meant to answer some specific questions about the details of stellar collisions and the resulting collision products. We can generalize our current results to all main-sequence star collisions in globular clusters under the limitations outlined in the previous papers. Therefore, we feel justified in claiming that no collision products have surface convection zones, and that collisions do not create long--lived circumstellar disks. The search to understand the angular momentum evolution of blue stragglers will have to turn to other avenues, possibly involving a more detailed look at the combination of stellar evolution and hydrodynamic instabilities after the collision." + }, + "0112/astro-ph0112417_arXiv.txt": { + "abstract": " ", + "introduction": "The cosmic microwave background (CMB) is now well known as a probe of the early universe. The temperature fluctuations in the CMB, especially the acoustic peaks in the angular power spectrum of CMB anisotropies, capture the physics of primordial baryon-photon fluid undergoing oscillations in the potential wells of dark matter \\cite{Huetal97}. In transit to us, CMB photons also encounter the large scale structure that defines the local universe; thus, several aspects of photon properties, such as the frequency or direction of propagation, are affected. In the reionized epoch, variations are imprinted when photons are Compton-scattered via electrons, moving with respect to the CMB. These gravitational and scattering effects, though some times insignificant compared to primary fluctuations, leave certain imprints in the anisotropy structure while inducing higher order correlations. These signatures can then be used to extract the properties of large scale structure that led to changes in the CMB temperature. Here, we will summarize what these signatures are, and, how they can be used to study the local universe with CMB photons. \\begin{figure}[t] \\begin{center} \\includegraphics[height=13pc,angle=-90]{cl.eps}\\includegraphics[height=13pc,angle=-90]{bl.eps}\\includegraphics[height=13pc,angle=-90]{convergence.eps} \\end{center} \\caption{{\\it Left}: Power spectrum for the temperature anisotropies in the fiducial $\\Lambda$CDM model with $\\tau=0.1$. The curves show the local universe contributions to CMB due to gravity (ISW and lensing) and scattering (Doppler, SZ effects, patchy reionization). {\\it Middle}: The lensing-secondary correlations and error bars show how well lensing-SZ correlation can be measured with Planck. {\\it Left}: CMB as a weak lensing experiment and error bars show the reconstruction of convergence with Planck.} \\label{fig:cl} \\end{figure} ", + "conclusions": "" + }, + "0112/hep-ph0112158_arXiv.txt": { + "abstract": "{We study the dynamics of active-sterile neutrino oscillations in the early universe using full momentum-dependent quantum-kinetic equations. These equations are too complicated to allow for an analytical treatment, and numerical solution is greatly complicated due to very pronounced and narrow structures in the momentum variable introduced by resonances. Here we introduce a novel dynamical discretization of the momentum variable which overcomes this problem. As a result we can follow the evolution of neutrino ensemble accurately well into the stable growing phase. Our results confirm the existence of a ``chaotic region\" of mixing parameters, for which the final sign of the asymmetry, and hence the SBBN prediction of $^4$He-abundance cannot be accurately determined.} ", + "introduction": "Recent observations from atmospheric and solar neutrinos from Super Kamiokande (SK)\\cite{SK1} and Sudbury Neutrino Observatory (SNO) \\cite{SNO} prefer conventional, mainly active-active mixings between ordinary electron, muon and tau neutrinos as the solutions to the observed flux deficits. However, while pure active-sterile solutions are disfavoured, sizable admixtures of sterile states in the observed fluxes are still allowed~\\cite{smirnov}. New sterile states are also needed in order to explain the LSND anomaly~\\cite{LosAl} by neutrino physics. Moreover, in a recent analysis of the so called ``2+2''-models\\footnote{Alternative ``3+1'' models fit less well with short baseline reactor experiments~\\cite{3p1schemes}.} for four-neutrino mixing, the solutions with nonzero sterile neutrino components were found to provide best global fits to the solar, atmospheric and reactor data~\\cite{stfluxes}. Corresponding limits on active-sterile mixings are quite generous, for example \\begin{eqnarray} \\sin^2\\theta_{\\mu s} &\\lesssim 0.48& \\qquad \\rm (Atmospheric) \\nonumber\\\\ \\sin^2\\theta_{e s} &\\lesssim 0.72& \\qquad \\rm (Solar \\;\\; LMA), \\label{labfluxes} \\end{eqnarray} where LMA refers to the Large Mixing Angle solution for the solar neutrino deficit. Active-sterile mixing, if realized, would have several interesting effects in astrophysical settings~\\cite{snova} and in particular for the evolution of the early universe \\cite{dol,ektdo,ekm,ektBig,ektat,ekmL,old,ftv,shi,eks1,fSc,Sorri,shifuNc}. In particular, sterile states could be brought into thermal equilibrium by mixing before nucleosynthesis, so that the resulting anomalous increase in the expansion rate of the universe would lead to overproduction of helium in disagreement with the observations~\\cite{dol,ektdo,ekm,ektBig,ektat,old}. Excluding the parameter sets leading to equilibration provides useful bounds on active-sterile mixing. Indeed, from results of ref.~\\cite{ektBig} one can infer that the sterile components in mass eigenstates responsible for atmospheric anomaly and LMA are constrained by \\begin{eqnarray} \\sin^2\\theta_{\\mu s} &\\lesssim 0.013& \\qquad \\rm (Atmospheric) \\nonumber\\\\ \\sin^2\\theta_{e s} &\\lesssim 0.026& \\qquad \\rm (Solar \\;\\; LMA) \\label{cosmofluxes} \\end{eqnarray} where we used a SBBN-limit of $N_\\nu \\lesssim 3.4$ for the number of effective neutrino degrees of freedom~\\cite{ektBig,recent}. These numbers correspond to the ``light\" case where the sterile component is the heavier of the mixing states; in the opposite, ``dark\" case, the corresponding limits are even stronger by one to two orders of magnitude. In either case the cosmological constraints are more stringent than those obtained in terrestrial laboratories. While quite generic, the cosmological bounds (\\ref{cosmofluxes}) do depend on some simple prior assumptions, the most important of which is that the primordial lepton asymmetry is not anomalously large\\cite{ekmL}. It may be possible to circumvent them in more complicated mixing schemes involving more sterile states (for a recent discussion, see~\\cite{dibSc}). In a particular attempt it was shown in ref.~\\cite{ftv} that for a large negative $\\delta m^2$ and a small enough mixing angle (so that the equilibration bounds of~\\cite{ektBig} can be avoided), say on the $\\nu_\\tau-\\nu_{s'}$-sector, resonant oscillations may trigger a rapid growth of the tau lepton asymmetry. This asymmetry could then become very large, $L_\\tau \\sim {\\cal O}(1)$ and change the natural prior condition of no significant initial lepton asymmetry for a mixing in any other sectors. Indeed, if created early enough, such an asymmetry could suppress $\\nu_\\mu-\\nu_s$-oscillations~\\cite{fvL1,ekmL}, obviating the bounds (\\ref{cosmofluxes}) for $\\nu_\\mu-\\nu_s$-mixing. However, as this scenario requires that the sterile state is the {\\em lighter} one (resonance condition), and that the mass splitting $|\\delta m^2|$ is large (to create $L_{\\tau}$ early enough), it is essentially excluded by the recent solar and atmospheric neutrinos combined with the direct constraint on the electron neutrino mass~\\cite{mainz}. Even with the mechanism of ref.~\\cite{ftv} by and large excluded, there can be important effects due to $L$-growth with smaller $|\\delta m^2|$. For example, a large homogeneous electron neutrino asymmetry $L_{\\nu_e}$ would directly influence the nucleosynthesis prediction for the helium-4 abundance through the reactions \\begin{eqnarray} n + e^+ &\\leftrightarrow & p + \\bar \\nu_e \\nonumber \\\\ p + e^- &\\leftrightarrow & n + \\nu_e \\;, \\end{eqnarray} which keep the neutron-to-proton ratio in equilibrium at early times. Changing the electron neutrino abundance could tilt the balance of these reactions with the possibility of either increasing $Y_{He}$ (negative $L_{\\nu_e}$), and hence strengthening the bounds (\\ref{cosmofluxes}), or decreasing $Y_{He}$ (positive $L_{\\nu_e}$), leading to weakening of the bounds (\\ref{cosmofluxes}). The physics involved with the resonant growth of the asymmetry is very complicated because of several vastly different physical scales both in the temporal direction and in momentum variable, which effect the evolution of the ensemble in an essential way. The relevant QKE's are nonlinear and strongly coupled internally via the asymmetry term, and therefore no useful analytical approximation exists for the problem\\footnote{ Some confusion in the field was caused by a recent analytical treatment~\\cite{dolHan}, whose results contradicted earlier numerical results~\\cite{shi,eks1,fSc}. Disagreement was eventually clarified in favour of numerical work in~\\cite{DiBetal}.}. Numerical solution is also greatly complicated by the various different physical scales, and while the phenomenon of asymmetry growth is fairly well established by now, the determinacy of the final {\\em sign} of the asymmetry is still not well understood. The ambiguity was first observed in~\\cite{shi}, and in~\\cite{eks1,Sorri} it was shown that this ``chaoticity\" occurs in certain well defined region of mixing parameters. References~\\cite{shi,eks1,Sorri} employed a numerical solution for the momentum averaged approximation of the QKE's however, and their results were challenged by ref.~\\cite{dolHan}, who claimed that the final asymmetry is always fully determined by the initial conditions. While the analysis of ref.~\\cite{dolHan} is now discredited~\\cite{DiBetal}, it still remains true that no numerically reliable momentum dependent analysis has so far shown the existence of the chaotic region.\\footnote{ Although ref.~\\cite{fdb} found support for chaoticity, their conclusion is not firm because of the reported loss of the numerical accuracy of their methods when approaching the potentially chaotic region.} It is important to settle the issue because, as mentioned above, large $L_{\\nu_e}$ could significantly alter the helium-4 abundance. Indeed, if the sign of $L_{\\nu_e}$ were found to be chaotic, then $Y_{He}$ could not be reliably determined from BBN~\\cite{eks1}. In this paper we study the dynamics of active-sterile neutrino oscillations in the early universe using full momentum-dependent quantum-kinetic equations (in homogeneous space). Our main technical improvement over the previous works is the introduction of a novel dynamical discretization method for the momentum variable which enables us to model accurately the density matrices over the entire momentum range, including the extremely pronounced and narrow structures close to resonant momenta. On physical context, our results confirm the existence of a \"chaotic\" region of mixing parameters, for which the final sign of the asymmetry is not deterministic. It also confirms the expectation~\\cite{eks1,fdb}, that the size of the chaotic region is somewhat smaller than indicated by the momentum averaged code~\\cite{eks1}. These results can be seen as giving more strength on the cosmological constraints on neutrino mixing parameters. The paper is organized as follows. In section 2 we set up the quantum kinetic equations for the problem. In section 3 we introduce the novel dynamically adjusted discretization of the momentum variable and in section 4 we set up our kinetic equations with this parametrization. The method we develop allows us to have enough resolution to solve the problem with relatively small number of momentum bins (we use at most 400 bins). In section 5 we display our numerical results for the key variables driving the oscillation and discuss the numerical stability of our solutions. Finally, in section 6 we present our results on the asymmetry growth and oscillations as well as the interpretation of the physical consequences, and the section 7 contains a summary and outlook. ", + "conclusions": "Active-sterile mixing, while constrained by the present solar and atmospheric data, are nevertheless required if one wishes to incorporate the LSND-neutrino anomaly into the neutrino models. Moreover, the active-sterile neutrino mixing would have a rich phenomenology in the early universe, where it provides a unique theoretical challenge in the form of a system for which the quantum effects play an essential role in the kinetic equations. Here we have studied the phenomenon of neutrino asymmetry growth in the early universe as a result of active sterile neutrino oscillations. Our main new contribution is the introduction of a novel method of discretizing the momentum variable such that the sharply pronounced structures near the resonances, which are here shown to drive the oscillation phenomena, can be treated numerically accurately. The only concession to rigor we made along the way to solution of the problem, was adding collision terms to regulate the sterile neutrino spectrum away from the resonances. However, we demonstrated that our regulators had no effect on the active neutrino evolution, and hence for the results presented here. Nevertheless, our treatment is obviously not adequate if the precise form of the sterile neutrino spectrum is important. This might be the case for the models where sterile neutrinos are invoked to provide a non-thermal component for the dark matter~\\cite{nonThDM}. We have demonstrated that even tiny changes in the oscillation parameters may drastically alter the oscillation pattern, and even change the sign of the final asymmetry. This behaviour does not occur for all oscillation parameters, but instead the dependence of the sign of $L$ on oscillation parameters is weak in the ``stable\" region (in general small $\\sin^22\\theta_0$ and large negative $\\delta m^2$) and strong in the chaotic region (in general large $\\sin^22\\theta_0$ and small negative $\\delta m^2$). These results are in qualitative agreement with our earlier findings, which were based on a momentum averaged treatment~\\cite{eks1}. Obviously, if active-sterile mixing were to be observed with parameters residing in the chaotic region, it would be, due to inevitable errors in the experimentally measured parameters, impossible to reliably determine the sign of the neutrino asymmetry created by that mixing in the early universe. Moreover, since large neutrino asymmetries (and electron neutrino asymmetry in particular) directly affect, in a $sign(L)$-dependent way, the weak interaction rates that determine proton-to-neutron ratio in the early universe, this indeterminacy would undermine our ability to accurately compute the BBN prediction for light element abundances for chaotic oscillation parameters. These conclusions hold in the spatially homogeneous calculation. However, the sensitivity of $L$-growth on initial conditions also lends to a speculation that in reality the initially inhomogeneous seed asymmetry might give rise to an inhomogeneous texture of domains of very large lepton asymmetries with oscillating sign. This phenomenon has been shown to occur in a one-dimensional, momentum averaged model~\\cite{eks2}, and could plausibly occur in a realistic three dimensional world. The implications of such a scenario for SBBN would obviously be very different, including the possibility of very efficient equilibration of sterile neutrinos via MSW-effect within the domain boundaries~\\cite{shiFuDo,eks2}. If this was the case, then the asymmetry growth mechanism would lead to even {\\em stronger} bounds on mixing than what is displayed in equations~(\\ref{cosmofluxes}). It is beyond the scope of this work (and the reach of the present computers) to study this phenomenon quantitatively in the momentum dependent case, however. Let us finally comment on the effects of large homogeneous lepton asymmetries for cosmic microwave background radiation (CMBR). The possibility of measuring $L$ by the Planck satellite data has been considered for example in~\\cite{Kinney:1999pd} and in~\\cite{CMBconstraints}. However, the only effect of $L$ on CMBR comes through the associated fluctuations in the energy density, and correspondingly only the total asymmetry has relevance. In the oscillation scenarios, such as discussed in this paper, the total asymmetry (here the sum of the sterile and active sector asymmetries) is conserved however, and hence the oscillation-induced asymmetries, in contrast to the situation with SBBN, would have no direct effect on CMBR." + }, + "0112/astro-ph0112001_arXiv.txt": { + "abstract": "We show that the dearth of brown dwarfs in short-period orbits around Solar-mass stars -- the brown dwarf desert -- can be understood as a consequence of inward migration within an evolving protoplanetary disc. Brown dwarf secondaries forming at the same time as the primary star have masses which are comparable to the initial mass of the protoplanetary disc. Subsequent disc evolution leads to inward migration, and destruction of the brown dwarf, via merger with the star. This is in contrast with massive planets, which avoid this fate by forming at a later epoch when the disc is close to being dispersed. Within this model, a brown dwarf desert arises because the mass at the hydrogen burning limit is coincidentally comparable to the initial disc mass for a Solar mass star. Brown dwarfs should be found in close binaries around very low mass stars, around other brown dwarfs, and around Solar-type stars during the earliest phases of star formation. ", + "introduction": "Recent surveys have demonstrated a high abundance of brown dwarfs in young clusters (Mart\\'\\i n et al. 2000), star forming regions (B\\'ejar et al. 2001), and the field (Kirkpatrick et al. 1999, 2000). Brown dwarfs are also reasonably common in wide binaries (Gizis et al. 2001). The sole well established exception is in close (semi-major axis $a \\ale 4 \\ {\\rm AU}$) binaries. The same radial velocity surveys that have been so successful in finding massive extrasolar planets show that brown dwarfs are rarely close binary companions to Solar-type stars (Marcy \\& Butler 2000; Halbwachs et al. 2000). This brown dwarf desert supports the conventional belief that massive planets and brown dwarfs form in distinctly different ways, and also suggests that some aspect of the formation or early evolution of brown dwarfs differs from that of stars. One possibility for explaining the desert is to postulate that the formation mechanism for brown dwarfs is entirely different from that of {\\em either} stars or giant planets. Reipurth \\& Clarke (2001), for example, suggest that brown dwarfs are objects whose growth towards stellar masses was curtailed by ejection from small multiple systems (see also Reipurth, Clarke \\& Delgado-Donate 2001). In their model, the desert arises because substellar objects that were {\\em not} ejected continued to accrete, eventually reaching stellar masses. In this paper, we investigate a less radical (and more limited) possibility. We show that the absence of brown dwarfs in close binaries with Solar-type stars could be due to orbital migration within an evolving protoplanetary disc. This is the same process that is invoked to explain the presence of massive planets at small orbital radii (Lin, Bodenheimer \\& Richardson 1996), and which may also play a role in the orbital evolution of cataclysmic variables (Taam \\& Spruit 2001) and supermassive black hole binaries (Gould \\& Rix 2000). The critical assumption is that brown dwarfs form contemporaneously with the star, and thus become embedded within a young and relatively massive protoplanetary disc. Under these conditions, we show that migration efficiently clears out the desert by forcing the brown dwarfs into mergers with the star. This is in contrast to the evolution of massive planets, which have a chance of escaping the same fate by forming later, when the disc is close to being dispersed and cannot drive migration through to merger (Trilling, Lunine \\& Benz 2001; Armitage et al. 2001). ", + "conclusions": "We have shown that inward orbital migration of brown dwarfs, within an evolving protoplanetary disc, can account for a low frequency of brown dwarfs as close binary companions to stars with masses $M_* \\sim M_\\odot$. Migration depletes the initial frequency of brown dwarf companions at all radii where there is significant viscous evolution of the protoplanetary disc. This region could extend out to several tens of AU. Brown dwarfs with initially larger orbital radii $a \\sim 10^2 \\ {\\rm AU}$ would be unaffected. In our specific model, migration leads to the destruction of all brown dwarfs with initial orbital radii $a \\ale 10 \\ {\\rm AU}$, via mergers with the star which cause significant stellar spin-up. Orbits with $0.1 \\ {\\rm AU} < a < 4 \\ {\\rm AU}$ are partially replenished by brown dwarfs migrating inwards from still greater radii, but the net brown dwarf frequency at these radii is still reduced by a factor of 5-10. The main prediction of the model is that brown dwarfs in close orbits ought to be up to an order of magnitude more common amongst the youngest pre-main-sequence stars (less than a Myr), as compared to the main sequence. For migration to occur, we also require that the initial disc mass be at least comparable to the mass of the brown dwarf. No brown dwarf desert is thus expected around the lowest mass hydrogen-burning stars ($0.1 - 0.2 \\ M_\\odot$), or indeed around other brown dwarfs, whose discs would have been too feeble to drive significant brown dwarf migration." + }, + "0112/hep-ph0112335_arXiv.txt": { + "abstract": "We discuss effects of cosmological moduli fields on the cosmic microwave background (CMB). If a modulus field $\\phi$ once dominates the universe, the CMB we observe today is from the decay of $\\phi$ and its anisotropy is affected by the primordial fluctuation in the amplitude of the modulus field. As a result, constraints on the inflaton potential from the CMB anisotropy can be relaxed. In addition, with the cosmological moduli fields, {\\sl correlated} mixture of adiabatic and isocurvature fluctuations may be generated, which results in enhanced CMB angular power spectrum at higher multipoles relative to that of lower ones. Such an enhancement can be an evidence of late time entropy production due to the cosmological moduli fields, and may be observed at on-going and future experiments. ", + "introduction": "In superstring theory \\cite{Polchinski}, it is well known that there are various flat directions parameterized by scalar fields. (Hereafter, we call these fields as ``moduli'' fields.) Since their potential is usually generated by effects of supersymmetry (SUSY) breaking, their masses are expected to be of the order of the gravitino mass. Although masses of the moduli fields can be as light as (or even lighter than) the electroweak scale, moduli fields do not affect collider experiments since their interactions are suppressed by inverse powers of the gravitational scale. Cosmologically, however, they may cause serious problems \\cite{PRL131-59}. If the mass of the moduli fields $m_\\phi\\mathop{}_{\\textstyle \\sim}^{\\textstyle <} O(10\\ {\\rm TeV})$, reheating temperature of the universe becomes lower than $\\sim 1\\ {\\rm MeV}$. With such a low reheating temperature, the success of the standard big-bang nucleosynthesis (BBN) is spoiled. For lighter moduli fields ($m_\\phi \\mathop{}_{\\textstyle \\sim}^{\\textstyle <} O(100\\ {\\rm MeV})$), they survive until today and overclose the universe. One solution to these difficulties is to push up the mass of the moduli fields \\cite{Heavyphi}. In particular, in Ref.\\ \\cite{NPB570-455}, it was pointed out that the scenario with heavy moduli can naturally fit into the framework of the anomaly-mediated SUSY breaking \\cite{AMSB}. Indeed the reheating temperature can be higher than $\\sim 1\\ {\\rm MeV}$ if $m_\\phi \\mathop{}_{\\textstyle \\sim}^{\\textstyle >}O(10\\ {\\rm TeV})$. In this case, the BBN occurs after the decay of the modulus field. Although the thermal history after the BBN is mostly the same as the standard one, cosmology before the modulus decay is completely different. Importantly, fluctuations of moduli fields affect the cosmic microwave background (CMB). In this talk, we consider this scenario and study its consequence in the CMB. ", + "conclusions": "We have studied the effects of the cosmological moduli fields on the CMB anisotropy. In the scenario with the cosmological moduli fields, {\\sl correlated} isocurvature fluctuation may exist in the baryonic sector which results in enhanced CMB angular power spectrum at higher multipoles relative to that at lower ones. In addition, even in the case where there is no isocurvature perturbation, the cosmological modulus field may have important implication to the model-building of the inflation." + }, + "0112/astro-ph0112093_arXiv.txt": { + "abstract": "We present the results of the first high angular resolution observations of SiO maser emission towards the star forming region W51-IRS2 made with the Very Large Array (VLA) and Very Long Baseline Array (VLBA). Our images of the H$_2$O maser emission in W51-IRS2 reveal two maser complexes bracketing the SiO maser source. One of these H$_2$O maser complexes appears to trace a bow shock whose opening angle is consistent with the opening angle observed in the distribution of SiO maser emission. A comparison of our H$_2$O maser image with an image constructed from data acquired 19 years earlier clearly shows the persistence and motion of this bow shock. The proper motions correspond to an outflow velocity of 80 km s$^{-1}$, which is consistent with the data of 19 years ago (that spanned 2 years). We have discovered a two-armed linear structure in the SiO maser emission on scales of $\\sim 25$ AU, and we find a velocity gradient on the order of $0.1$ km s$^{-1}$ AU$^{-1}$ along the arms. We propose that the SiO maser source traces the limbs of an accelerating bipolar outflow close to an obscured protostar. We estimate that the outflow makes an angle of $< 20^{\\circ}$ with respect to the plane of the sky. Our measurement of the acceleration is consistent with a reported drift in the line-of-sight velocity of the W51 SiO maser source. ", + "introduction": "Since the discovery of the first bipolar outflows in star-forming regions \\citep{SNELL+80}, significant progress has been made in understanding the large-scale characteristics of such outflows \\citep[{e.g.,}][]{BACHILLER96}. However, the large columns of gas and dust that obscure massive protostars hinder traditional optical or infra-red observations of at least the inner $\\sim 100$ AU of these outflows, where the exciting protostars reside. This has made it difficult to obtain sufficient data to understand well the process of high-mass star formation. Radio frequency maser emission, which traces velocity coherent clumps of gas entrained in large-scale bulk mass motions, is unattenuated by the neutral gas and dust around massive protostars. Moreover, because masers are compact, high-brightness sources, high angular resolution radio interferometry can be used to probe the structure and kinematics of these outflows on angular scales of $\\lesssim 1$ AU for many Galactic sources \\citep[{e.g.,}][]{GREENHILL+98, PATEL+00}. Many protostellar outflows exhibit H$_2$O maser emission at 22 GHz \\citep[{e.g.,}][]{HENNING+92,FELLI+92} which traces shocks in dust-laden gas close to the exciting protostars \\citep{ELITZUR92}. SiO masers, which are a common feature of late-type stars, are known to occur in only three regions of star formation: W51-IRS2, Sgr-B2 MD5, and Orion-KL \\citep{HASEGAWA+86, SB74}. Maser action in vibrationally excited states of SiO at 43 GHz requires higher temperatures ($>10^3$ K) and is more closely associated with exciting sources than is H$_2$O maser emission \\citep{ELITZUR92}. The only non-stellar SiO maser source that has been well studied is the one in Orion-KL, where the maser emission traces an outflow within $\\sim 100$ AU of an obscured massive protostar \\citep{GREENHILL+98, DLP99}. W51-IRS2 is an embedded infrared source ($L_{\\rm tot} \\sim 2.8 \\times 10^6 \\: L_{\\odot}$; Erickson \\& Tokunaga 1980) in the well-known high-mass star forming region W51. We adopt a distance of 7 kpc to W51 based on maser proper motion studies \\citep{GENZEL+81}. The region around IRS2 contains an edge-brightened cometary HII region called W51d \\citep{MARTIN72, WC89, GJW93} that is associated with a 2.2 $\\mu$m point source \\citep{GW94} and a peak in 12 $\\mu$m emission \\citep{OKAMOTO+01}. An ultra-compact HII region called W51d2 \\citep{MARTIN72} is associated with NH$_3$ and methanol masers \\citep{GJW93, MCB01} as well as a feature in the 12 $\\mu$m continuum \\citep{OKAMOTO+01}. W51-IRS2 also contains an H$_2$O maser complex called W51 North, for which detailed distributions and proper motions have been observed \\citep{SCHNEPS+81}, OH maser emission \\citep{GM87}, and an unresolved SiO maser source \\citep{HASEGAWA+86, UKITA+87, MORITA+92}. The strongest H$_2$O masers, the OH masers, and the SiO masers, are found in a compact region of W51 North termed the ``Dominant Center'' by Schneps et al. (1981). Although no infrared or radio continuum sources have been detected within $\\sim 2''$ of it \\citep{GJW93, WC89, OKAMOTO+01}, the Dominant Center does coincide with thermal emission from several molecular species including NH$_3$(1,1), NH$_3$(2,2), NH$_3$(3,3), CS, and CH$_3$CN \\citep{ZH97,ZHO98}. The spectrum of the H$_2$O maser emission covers $V_{\\rm LSR} = -30$ to 130 km s$^{-1}$, and is peaked around 60 km s$^{-1}$ \\citep{SCHNEPS+81}. The SiO emission is peaked around 45 km s$^{-1}$ and covers $V_{\\rm LSR} = 40$ to 50 km s$^{-1}$, lying within the velocity range of the H$_2$O maser emission \\citep{MORITA+92}. The spectra for various molecular species are peaked around $V_{\\rm LSR} \\sim 60$ km s$^{-1}$, with typical linewidths of $\\sim 20$ km s$^{-1}$ \\citep{ZH97,ZHO98}. In this paper, we present the first high angular resolution observations of the SiO maser emission in W51-IRS2. In \\S \\ref{sec:obs} and \\S \\ref{sec:results}, we describe interferometric observations of the SiO and H$_2$O maser emission. We interpret our results in the context of a bipolar outflow model in \\S \\ref{sec:dc-disc}. We also show that previous measurements of the proper motions of H$_2$O masers and the velocity drifts of SiO masers provide support for our model. ", + "conclusions": "\\label{conclns} We have resolved the structure of the W51-IRS2 SiO maser source, and linked it to a protostellar outflow associated with two long-known sites of intense H$_2$O maser emission. We propose that the masers trace an accelerating bipolar protostellar outflow inclined $<20^{\\circ}$ with respect to the plane of the sky, which may detectably rotate about the axis of flow within tens of AU of the protostar. We estimate the position angle of the flow to be $\\sim 105^{\\circ}$ up to 4200 AU from the central star. The proper motions of H$_2$O maser clusters bracketing the SiO maser source indicate an outflow velocity of $\\sim 80$ km s$^{-1}$ along this position angle, and one of these clusters appears to trace a bow shock that subtends an angle consistent with the opening angle suggested by the two limbs of SiO maser emission. We estimate the acceleration of the outflow to be $\\sim 0.5$ km s$^{-1}$ yr$^{-1}$, which is consistent with the 0.4 km s$^{-1}$ yr$^{-1}$ line-of-sight velocity drift measured by Fuente et al. (1989). In the larger context, this bipolar flow lies in within a $\\sim 10^4$ AU cloud core hosting multiple centers of H$_2$O maser emission. The outflow does not extend to the limits of the core or to the other centers of H$_2$O maser emission. We estimate the line-of-sight velocity of the protostar that is driving the bipolar outflow to be $\\sim 47$ km s$^{-1}$, significantly different than the systemic velocity of the NH$_3$ core (60 km s$^{-1}$). We suggest that there may be several star-forming fragments within this core, perhaps marked by the other centers of H$_2$O maser activity." + }, + "0112/astro-ph0112570_arXiv.txt": { + "abstract": "In view of the extensive evidence of a tight inter-relationship between spher\\-oidal galaxies (and galactic bulges) and massive black holes hosted at their centers, a consistent model must deal jointly with the evolution of the two components. We describe one viable model, which successfully accounts for the local luminosity function of spheroidal galaxies, their photometric and chemical properties, deep galaxy counts in different wavebands, including those in the (sub)-mm region which proved to be critical for current semi-analytic models stemming from the standard hierarchical clustering picture, clustering properties of SCUBA galaxies, of EROs, and of LBGs, as well as for the local mass function of massive black holes and for quasar evolution. Predictions that can be tested by surveys carried out by SIRTF are presented. ", + "introduction": "The hierarchical clustering model with a scale invariant spectrum of density perturbations in a Cold Dark Matter (CDM) dominated universe has proven to be remarkably successful in matching the observed large-scale structure as well as a broad variety of properties of galaxies of different morphological types (e.g. \\cite{Cole2000,Granato2000}). However, serious shortcomings of this scenario have also become evident in recent years. The critical point can be traced back to the relatively large amount of power on small scales predicted by this model which would imply far more dwarf galaxies or substructure clumps within galactic and cluster mass halos than are observed (the so-called ``small-scale crisis'' (\\cite{Haiman2001,Somerville2001,KamionkowskiLiddle2000,Mooreetal1999}), unless star formation in small objects is strongly suppressed (or the small scale power is reduced by modifying the standard model). At the other extreme of the galaxy mass function we have another strong discrepancy with model predictions, that we might call ``the massive galaxy crisis'': even the best semi-analytic models (\\cite{Granato2000,Devriendt2000}) hinging upon the standard picture for structure formation in the framework of the hierarchical clustering paradigm, fall short by a substantial factor (up to about 10) to account for the (sub)-mm (SCUBA and MAMBO) counts of galaxies, most of which are probably massive objects undergoing a very intense star-burst (with star formation rates $\\sim 1000\\,\\hbox{M}_\\odot\\,\\hbox{yr}^{-1}$) at $z>2$ (see, e.g. \\cite{Dunlop2001}). Recent optical data confirm that most massive ellipticals were already in place and (almost) passively evolving up to $z\\simeq 1$--1.5, implying that they were fully assembled by $z \\sim 2.5$, although the issue is still somewhat controversial (\\cite{RenziniCimatti1999,Ferguson2000,Daddi2000b,Martini2001,Rodighiero2001,Cohen2001,Im2001,McCarthy2001}). These data are more consistent with the traditional ``monolithic'' approach whereby giant ellipticals formed most of their stars in a single gigantic starburst at substantial redshifts, and underwent essentially passive evolution thereafter. On the contrary, in the canonical hierarchical clustering paradigm the smallest objects collapse first and most star formation occurs, at relatively low rates, within relatively small proto-galaxies, that later merged to form larger galaxies. Thus, the expected number of galaxies with very intense star formation is far less than detected in SCUBA and MAMBO surveys and the surface density of massive evolved ellipticals at $z\\gsim 1$ is also smaller than observed. The ``monolithic'' approach, however, is inadequate to the extent that it cannot be fitted in a consistent scenario for structure formation from primordial density fluctuations. ", + "conclusions": "" + }, + "0112/astro-ph0112436_arXiv.txt": { + "abstract": "\\xmm\\ \\rgs\\ spectra of \\mcg\\ and \\mrk\\ exhibit complex discrete structure, which was interpreted in a paper by \\citet{grazie01} as evidence for the existence of relativistically broadened Lyman alpha emission from carbon, nitrogen, and oxygen, produced in the inner-most regions of an accretion disk around a Kerr black hole. This suggestion was subsequently criticized in a paper by \\citet{lee01}, who argued that for \\mcg, the \\chandra\\ \\hetg\\ spectrum, which is partially overlapping the \\rgs\\ in spectral coverage, is adequately fit by a dusty warm absorber model, with no relativistic line emission. We present a reanalysis of the original \\rgs\\ data sets in terms of the \\citet{lee01} model. Specifically, we show that: (1) The explicit model given by \\citet{lee01} differs markedly from the \\rgs\\ data, especially at longer wavelengths, beyond the region sampled by the \\hetg; (2) Generalizations of the \\citet{lee01} model, with all parameters left free, do provide qualitatively better fits to the \\rgs\\ data, but are still incompatible with the detailed spectral structure; (3) The ionized oxygen absorption line equivalent widths are well-measured with the \\rgs\\ for both sources, and place very tight constraints on both the column densities and turbulent velocity widths of \\ion{O}{7} and \\ion{O}{8}. The derived column densities are well below those posited by \\citet{lee01}, and are insufficient to play any role in explaining the observed edge-like feature near $17.5$~\\AA; (4) The lack of a significant neutral oxygen edge near $23$~\\AA\\ places very strong limits on any possible contribution of absorption to the observed structure by dust embedded in a warm medium; (5) The original relativistic line model with warm absorption proposed by \\citet{grazie01} provides a superior fit to the \\rgs\\ data, both in the overall shape of the spectrum and in the discrete absorption lines. We also discuss a possible theoretical interpretation for the putative relativistic Lyman alpha line emission in terms of the photoionized surface layers of the inner regions of an accretion disk. While there are still a number of outstanding theoretical questions about the viability of such a model, it is interesting to note that simple estimates of key parameters are roughly compatible with those derived from the observed spectra. ", + "introduction": "High-resolution soft X-ray spectra of Seyfert~1 galaxies acquired with the grating spectrometers onboard the \\chandra\\ and \\xmm\\ observatories have allowed us, for the first time, to investigate both the dynamics of the extended absorbing media in these systems, and the structure of the intrinsic spectra that originate in the vicinity of the central massive black holes. To date, narrow X-ray absorption lines have been detected in many objects, including NGC~5548 \\citep{kaastra00}, NGC~3783 \\citep{kaspi00,kaspi01}, \\mcg\\ and \\mrk\\ \\citep{grazie01,lee01}, IRAS~13349+2438 \\citep{sako01}, NGC~4051 \\citep{collinge01}, and Mrk~509 \\citep{pounds01}. On the other hand, the spectra of several sources show no evidence for absorption, but, instead, exhibit complex spectral features and temporal behavior that cannot be explained by simple continuum emission models, and are more likely to be related to the activity of the central engine (e.g., Ton~S180, \\citealt{turnertj01}; NGC~4593, \\citealt{katrien01}). In an earlier paper presenting high-resolution spectra of \\mcg\\ and \\mrk\\ obtained with the Reflection Grating Spectrometer (\\rgs) onboard \\xmm\\ \\citep{grazie01}, we suggested that the soft X-ray emission of these two objects in the $\\lambda = 5 - 38$~\\AA\\ ($E = 0.35 - 2.5 ~\\rm{\\kev}$) band includes significant contributions from discrete emission features produced through X-ray illumination of an accretion disk around the central black hole. The observed spectral structure matches well with what is expected for relativistically broadened recombination lines of hydrogen-like carbon, nitrogen, and oxygen originating from the inner regions of the disk in a Kerr metric. That interpretation significantly challenges conventional models of the global soft X-ray spectra of \\agn, where it is usually assumed that all discrete structure in the soft X-ray band is due to absorption features produced in an extended medium far away from the central black hole. In a subsequent paper, \\citet{lee01} questioned our interpretation based on their analysis of a non-simultaneous \\chandra\\ \\hetg\\ observation of \\mcg. They claimed that the $\\lambda < 25$~\\AA\\ ($E > 0.5 ~\\rm{\\kev}$) region can be well-reproduced solely by a dusty warm absorber model superimposed on a smooth continuum, and does not require the presence of any relativistic disk emission lines. They interpret the excess soft X-ray emission above the high energy power law continuum as either a blackbody component, or a steepening of the power law in the soft X-ray band. The purpose of this paper is to provide a reanalysis of the \\xmm\\ data presented by \\citet{grazie01}, with more detailed attention to the properties of the discrete absorption components in both sources. For \\mcg, we explicitly test the spectral model proposed by \\citet{lee01}, and demonstrate that, contrary to their claim, it cannot account for all of the features observed in the \\rgs\\ spectrum. In particular, the absorption model fails to self-consistently fit the details of both the oxygen line absorption and the apparent edge-like structure present in the data. Further, the \\citet{lee01} hypothesis that much of the inferred absorption can be ascribed to warm dust, is clearly ruled out by the \\rgs\\ spectra. We reconfirm our earlier conclusion that the global spectrum strongly favors the intrinsic nuclear radiation to be dominated by discrete emission line features that are expected to be formed in an X-ray photoionized accretion disk. The observed line profiles indicate that they are distorted through relativistic beaming and smeared by strong gravitational effects in the vicinity of the black hole. This paper is organized as follows. In \\S2, we summarize the observations and describe the procedures adopted for data reduction. While we focus primarily on the spectra obtained by the \\rgs, we adopt continuum parameters derived from the \\epic\\ data in the $2 - 10 ~\\rm{\\kev}$ region, which is mostly outside the \\rgs\\ bandpass. The procedures and results are summarized in \\S3. In \\S4, we discuss the model proposed by \\citet{lee01} and compare it directly to the \\rgs\\ data. We also describe in detail our measurements of the ionized absorber parameters, as well as the possible existence of absorption by dust and the limits on the column densities. Having characterized the contributions of the extended absorbing medium, we demonstrate that the intrinsic spectra of both \\mcg\\ and \\mrk\\ are not smooth, but instead consist of discrete jumps. We interpret these jumps as blue emission edges of carbon, nitrogen, and oxygen Ly$\\alpha$ recombination lines produced in the inner regions of a relativistic accretion disk, which are described in \\S5. We briefly summarize our results in \\S6. ", + "conclusions": "By taking advantage of the extended wavelength coverage and high statistical quality provided by \\xmm, we have shown that the dusty warm absorber model posited by \\citet{lee01} cannot explain the measured \\rgs\\ spectra of either \\mcg\\ or \\mrk. The ionized oxygen column densities and turbulent velocity widths are well-constrained by the data, and eliminate the possibility that ionized oxygen makes any significant contribution to the observed edge-like feature near 17.5~\\AA. The absence of any discernible neutral oxygen edge, and the detailed structure of the 17.5~\\AA\\ feature further suggest that absorption by iron oxides associated with dust is negligible in both spectra. By contrast, the model involving relativistic Ly$\\alpha$ emission lines of carbon, nitrogen, and oxygen, with some contribution from an ionized absorber, quantitatively provides a superior fit. The parameters characterizing these emission lines are roughly compatible with what would be expected for the surface layers of an irradiated accretion disk. While there are still outstanding theoretical questions associated with this interpretation, the relativistic line model more successfully reproduces the spectra of these two objects." + }, + "0112/astro-ph0112350_arXiv.txt": { + "abstract": "By directly probing mass distributions, gravitational lensing offers several new tests of the CDM paradigm. Lens statistics place upper limits on the dark matter content of elliptical galaxies. Galaxies built from CDM mass distributions are too concentrated to satisfy these limits, so lensing extends the ``concentration problem'' in CDM to elliptical galaxies. The central densities of the model galaxies are too low on $\\sim\\!10$ pc scales to agree with the lack of central images in observed lenses. The flux ratios of four-image lenses imply a substantial population of dark matter clumps with a typical mass $\\sim\\!10^{6}\\ M_\\odot$. Thus, lensing implies the need for a mechanism that reduces dark matter densities on kiloparsec scales without erasing structure on smaller scales. ", + "introduction": "The popular Cold Dark Matter (CDM) paradigm is facing several challenges on small scales (e.g., \\cite{Moore01}). The dynamics of spiral galaxies, especially rotation curves and fast-rotating bars, suggest that in observed galaxies dark matter halos are much less concentrated than predicted by CDM (e.g., \\cite{Debattista}, \\cite{deBlok}), although this conclusion is still controversial (e.g., \\cite{vandenBosch}). The number of satellite dwarf galaxies in the Local Group is much smaller than the number of subhalos in CDM simulations \\cite{Klypin}, \\cite{Moore99}, although the discrepancy may be explained by the astrophysics of star formation rather than by the physics of the dark matter particle \\cite{Bullock00}. These tests of CDM are limited, however, by uncertainties in interpreting luminous tracers of the potential. Gravitational lensing offers a different test that probes mass distributions directly. Strong lensing by galaxies robustly determines the total mass in the inner 5--10 kpc of lens galaxies, which are predominantly elliptical galaxies. It also offers the possibility to detect small-scale mass concentrations in galaxy halos \\cite{Chiba}, \\cite{Dalal}, \\cite{Keeton2}, \\cite{Mao}, \\cite{Metcalf}. Lensing thus offers new tests of CDM that avoid dynamical uncertainties and extend the tests from spiral galaxies to ellipticals. ", + "conclusions": "Lens statistics imply that the dark matter densities in the inner parts of elliptical galaxies are lower than predicted by CDM, in agreement with the conclusion from dynamical analyses of spiral galaxies. The CDM paradigm must therefore be modified to reduce dark matter densities on kiloparsec scales. Various mechanisms have been proposed ranging from astrophysics (disk bars that erase dark matter cusps \\cite{Weinberg}) to cosmology (a tilted power spectrum \\cite{Alam}) to particle physics (dark matter that is not collisionless and cold \\cite{Bode}, \\cite{Colin}, \\cite{Spergel}). Lensing also implies that lens galaxies have high densities on small scales ($\\la\\!10$ pc). The central densities of galaxies must be much higher than predicted in CDM model galaxies to explain the absence of central or ``odd'' images in observed lenses. The flux ratios in four-images lenses imply that a substantial fraction of the dark matter ($\\sim\\!2\\%$) lies in small-scale clumps rather than a smooth halo component \\cite{Dalal}, and B1422+231 suggests that a typical clump mass is $\\sim\\!10^{6}\\ M_\\odot$ \\cite{Keeton2}. Thus, while lensing supports other evidence that a mechanism is needed to reduce dark matter densities on kiloparcsec scales, it also suggests that the mechanism must {\\it not\\/} remove structure on small scales --- which argues against changing the nature of the dark matter particle. \\vskip\\baselineskip {\\bf Acknowledgments.} Support for this work was provided by Steward Observatory at the University of Arizona, and by NASA through Hubble Fellowship grant HST-HF-01141.01-A from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS5-26555." + }, + "0112/astro-ph0112399_arXiv.txt": { + "abstract": "Earth is unusual in bearing life, and in having a large moon. A number of authors have suggested a possible connection between the two, e.g. through lunar stabilisation of the earth's obliquity, or through the effects of the oceanic tides. The various suggestions are reviewed. ", + "introduction": "The properties of the universe that we observe must be consistent with the evolution of carbon-based life within it. This observational selection effect is known as the {\\it weak anthropic principle} (Barrow \\& Tipler 1986). It has been invoked to explain a number of otherwise unlikely coincidences such as the nuclear resonance that allows carbon to form in stellar interiors, the big numbers coincidence (ratio of strengths of electromagnetic and gravitational forces $\\sim$ 10$^{40}$ $\\sim$ current size of the observable universe in proton diameters), and, more recently, the smallness of the cosmological constant (Efstathiou 1995) and the amplitude of the primordial density fluctuations which seeded the growth of galaxies and clusters of galaxies (Tegmark \\& Rees 1998). A similar selection effect will apply in our local astrophysical environment. Clearly, the luminosity and lifetime of the sun, and the shape and size of the earth's orbit, must be such as to maintain the earth's surface, for a long period, at a temperature suitable for the evolution of organic life (e.g. Kasting et al 1993). However, our solar system may also be atypical in other respects: \\newline $\\bullet$ The sun's metallicity is unusually high for its age (Whittet 1997, Gonzalez 1999), perhaps reflecting the higher probability of planetary systems being associated with high-metallicity parent stars. \\newline $\\bullet$ The sun's luminosity may be unusually stable (Gonzalez 1999). \\newline $\\bullet$ Jupiter's role in ejecting comets from the solar system could have been crucial in protecting the young earth from life-inhibiting impacts (Wetherill 1995). Solar systems with Jupiter-like planets at similar radii may thus not be typical. \\newline $\\bullet$ The sun may be orbiting the galaxy close to the co-rotation circle (Mishurov \\& Zenina 1999), which minimises the number of spiral-arm crossings, and consequent disruption of the solar system (e.g. by nearby supernovae, tidal effects). Although it's possible that none of the above features of our solar system is {\\it essential} to the evolution of life on earth, the probability of our observing them is enhanced if they increase the probability that intelligent life will develop, i.e. it would not be surprising to observe any feature $F$ if its a priori probability, $p_F$, satisfies: \\hspace*{15mm} $p_{F}>p_{life}(0)/p_{life}(F)$ \\hfill (1) where $p_{life}(0)$ and $p_{life}(F)$ are the probabilities of intelligent life evolving respectively in the absence and in the presence of feature $F$. Several authors (e.g. Butler 1980, Comins 1993) have suggested that there might be a link between our large moon, arguably an a priori unlikely feature of the earth's environment, and the evolution of life on earth, i.e. that: \\hspace*{15mm} $p_{moon}>p_{life}(0)/p_{life}(moon)$ \\hfill (2) All 3 parameters in this inequality are unknown. Below we consider what is known about the origin of the moon, and the origin of life, and how the former might affect the latter. ", + "conclusions": "In studies of the solar system, as in cosmology, we are dealing with a unique example (so far), and we must beware the effects on our observations of anthropic selection. The possibility that the presence of the moon has affected the origin or evolution of life through one of the mechanisms noted above implies that: \\newline (1) hypotheses about the origin of the moon cannot be judged solely on the basis of a priori likelihood (they need only satisfy equation 2); \\newline (2) large moons may be useful pointers in the search for life-bearing planets. More generally, caution must be exercised in interpreting, and generalising from, unusual features of our local astrophysical environment; they may turn out to be anthropically selected, and atypical. Indeed, the earth and its environment may be very special (Ward \\& Brownlee 2000), and this might explain the puzzling lack of evidence for intelligent life elsewhere in the universe (Tipler 1980, Wesson 1990). \\theendnotes" + }, + "0112/astro-ph0112485_arXiv.txt": { + "abstract": "{ We report on the results of a series of X--ray observations of the transient black hole candidate XTE J2012+381 during the 1998 outburst performed with the BeppoSAX satellite. The observed broad-band energy spectrum can be described with the superposition of an absorbed disk black body, an iron line plus a high energy component, modelled with either a power law or a Comptonisation tail. The source showed pronounced spectral variability between our five observations. While the soft component in the spectrum remained almost unchanged throughout our campaign, we detected a hard spectral tail which extended to 200 keV in the first two observations, but became barely detectable up to 50 keV in the following two. A further re-hardening is observed in the final observation. The transition from a hard to a soft and then back to a hard state occurred around an unabsorbed 0.1--200 keV luminosity of $10^{38}\\ergs$ (at 10 kpc). This indicates that state transitions in XTE 2012+281 are probably not driven only by mass accretion rate, but additional physical parameters must play a role in the evolution of the outburst. } ", + "introduction": "The transient X--ray source XTE J2012+381 was discovered with the Rossi X--ray Timing Explorer All Sky Monitor (RXTE-ASM) on May 24, 1998 at a level of 23 mCrab (2--12 keV; Remillard, Levine \\& Wood 1998; see Fig. \\ref{lc}). Within 3 days the source raised to an average level of 88 mCrab (May 27, 1998), reaching values as high as 110 mCrab (2--10 keV, RXTE Proportional Counter Array - PCA; Marshall \\& Strohmayer 1998). An ASCA observation on May 29, 1998 (MJD 50962) show\\-ed the source at a level 150 mCrab (2--10 keV; White et al. 1998). The ASCA Gas Imaging Spectrometer 0.5--10 keV spectrum could be well described by the superposition of a multicolor disk black body and a power law model. The temperature of the innermost disk radius was $k\\,T=0.76\\pm 0.01$ keV, the power law photon index $\\Gamma=2.9\\pm0.1$ and the absorption column density $N_H=(1.29\\pm0.03)\\times 10^{22}\\cmdue$ (White et al. 1998). The presence of a soft thermal component (equivalent temperature $\\sim 1$ keV) plus a hard power law is considered to be a characteristic feature of black hole candidates (e.g. Tanaka \\& Shibazaki 1996). The outburst evolution as observed by the RXTE-ASM is shown in Fig. \\ref{lc}. An extensive analysis of the available RXTE PCA pointings was carried out by Vasiliev, Trudolyubov \\& Revnivtsev (2000). Very Large Array (VLA) observations on 31 May, 1998 led to the identification of a radio source (Hjellming, Rupen \\& Mioduszewski 1998a). The radio source was located at R.A.=20$^{\\rm h}$12$^{\\rm m}$37$^{\\rm s}.67$ and Dec.=+38$\\deg$11$^{\\rm '}$01$^{\\rm ''}.2$ (equinox 2000) within the $\\sim 1{\\rm'}$ RXTE $90\\%$ error circle (Hjellming, Rupen \\& Mioduszewski 1998b). Despite the large column density (converting to a V-band extinction of about 7 mag) inferred from the X--ray spectrum, the large flux variations characteristic of X--ray novae made possible the identification of a faint red (R=20.1; V=21.3) counterpart during the X--ray source outburst, at a position consistent with the location of the radio counterpart (Hynes \\& Roche 1998; Hynes et al. 1999). Here we report on an observation campaign carried out with the Italian-Dutch satellite BeppoSAX (Boella et al. 1997a), aimed at studying in detail the outburst of XTE J2012+381. Five observations were carried out starting from May 28, 1998 through July 8, 1998. The X--ray spectra and light curves are discussed in Section 2. Section 3 is dedicated to the discussion of the results; our conclusions are summarised in Section 4. \\begin{figure*}[htb] \\psfig{file=x2012.ps,height=12 truecm} \\caption[h]{X--ray light curve of XTE J2012+381 as observed with the RXTE-ASM in the three energy band (1.3--3, 3--5 and 5--12 keV). The arrows in the upper panel mark the epochs of the five BeppoSAX observations; the arrow in the middle panel marks the epoch of the ASCA observation whereas the bars at the bottom of the lower panel show the epochs of the RXTE pointed observations.} \\label{lc} \\end{figure*} \\begin{table*} \\caption{Summary of BeppoSAX observations.} \\begin{tabular}{cc|cc|cc|cc|cc} Obs.& Start & LECS & LECS & MECS & MECS & HPGSPC & HPGSPC & PDS & PDS \\\\ num.& time & Exp. time & Count rate & Exp. time & Count rate & Exp. time & Count rate& Exp. time & Count rate \\\\ & MJD & (s) & c s$^{-1}$ & (s) & c s$^{-1}$ & (s) & c s$^{-1}$& (s) & c s$^{-1}$ \\\\ \\hline 1 &50961.554& 11540 & 24.4 & 23641 & 19.6 & 11544\t & 11.2 & 11973 & \\ 2.7 \\\\ 2 &50974.535& \\ 4151 & 20.1 & \\ 7916 & 17.5 & \\ 2769 &\\ 1.4 & \\ 2396 & \\ 1.4 \\\\ 3 &50982.365& 16495 & 22.9 & 27339 & 20.0 & 11160 &\\ 0.6 & 10281 & \\ 0.2 \\\\ 4 &50987.250& 19388 & 26.7 & 31500 & 22.2 & -- & -- & 13282 & $<0.2^*$ \\\\ 5 &51002.854& 15214 & 19.7 & 27142 & 17.0 & -- & -- & 12037 & 0.3 \\\\ \\end{tabular} \\noindent Count rates are in the energy bands used in the spectral analysis (see text). \\noindent $^*$ $3\\,\\sigma$ upper limit. \\label{date} \\end{table*} ", + "conclusions": "BeppoSAX observations that covered the 1998 outburst of XTE J2012+381 found several similarities with known BHCs. In particular, the X--ray spectrum could be well described by an ultrasoft component (disk black body) with the addition of a variable hard power law. One can derive an upper limit on the black hole mass by interpreting the inner disk radius from the spectral fits as the last marginally stable orbit. The lowest values are obtained for the first observation with an equivalent radius $R=33$ km. For a disk inclination of $0\\deg$, we derive $M\\gsim 3.7\\,d_{10}\\msole$ ($M\\gsim 22\\,d_{10}\\msole$) for a non-rotating (maximally rotating) black hole. Notice that the absolute value inferred for the disk radius, due to the approximations in the disk black body model used, is to be considered also a lower limit (Merloni, Fabian \\& Ross 2000). One of the main results of the present paper is the different spectral states observed during the outburst evolution of XTE 2012+281. The column density, the parameters of the soft component and iron line remained almost constant. On the other hand the hard tail changed considerably across different observations. During the first flare (between MJD 50960 and 50982) the X--ray spectrum was hard with a power law index $\\Gamma=2.2\\pm0.1$. In the second flare (between MJD 50982 and 51002) the X--ray spectrum firstly became soft and then it re-hardened (between MJD 51003 and 51014). At the outburst end, there is a hint for a further softening. In passing, we note that BeppoSAX and RXTE spectral fits are in good agreement, despite the poorer resolution of RXTE spectra. Moreover, our timing analysis put in evidence a higher r.m.s. variability in the interval during which the spectrum of XTE 2012+281 is soft. The transition from a hard to a soft and then back to a hard state occurs around the same count rate in the RXTE-ASM around ($\\sim 9$ c s$^{-1}$). Clearly for such an absorbed source the ASM cannot provide a fair estimate of the source luminosity. Comparing BeppoSAX (unabsorbed) luminosities derived from the spectral decomposition we see that this threshold is around $10^{-8}\\ergs\\cmdue$, i.e. $10^{38}\\,d^2_{\\rm 10 kpc}\\ergs$. Above this flux level, during the first flare the spectrum of XTE 2012+281 is hard, whereas during the second flare above the same flux level the source spectrum is soft. Simultaneously, the r.m.s. fractional variability increases only slightly from $\\sim 2\\%$ to $\\sim 3\\%$. This likely indicate that state transitions in XTE 2012+281 are not driven only by mass accretion rate and at least one other parameter is needed. In particular, the soft component remains almost stable during the BeppoSAX and RXTE observational campaigns (within a factor of a few) whereas the hard component varied considerably. In this regard we note that spectral parameters such as the plasma temperature $T_{\\rm C}$ or the high energy cut-off vary considerably across the transition strengthening the idea that the second parameter involved in state transitions is related to a hot corona." + }, + "0112/astro-ph0112166_arXiv.txt": { + "abstract": "The Parkes multibeam pulsar survey is covering a $10\\deg$-wide strip of the southern Galactic plane from $l=260\\deg$ to $l=50\\deg$. It utilizes a 13-beam receiver operating in the 20-cm band on the Parkes 64-m radio telescope and is much more sensitive than any previous large-scale survey. Most of the 608 pulsars discovered so far are relatively distant and many are young, with 37 having a characteristic age of less than $10^5$ years. At least one of these is associated with a supernova remnant and four other probable associations are suggested. Several multibeam pulsars have high values of the parameter $\\dot E/d^2$ and are within the position error contours of unidentified EGRET gamma-ray sources. These possible associations will be tested with the advent of new gamma-ray telescopes. ", + "introduction": "The Parkes multibeam pulsar survey is a large-scale survey of the southern Galactic plane from $l=260\\deg$ to $l=50\\deg$ and with $|b|<5\\deg$ using the Parkes 64-m radio telescope and the multibeam receiver. This receiver has 13 beams arranged in a hexagonal pattern, each with dual linear polarization and a bandwidth of 288 MHz centered on 1374 MHz. The average system temperature is about 21~K, corresponding to a system-equivalent flux density of about 30 Jy. Signals from each polarization of each beam are filtered to give 96 3-MHz channels, the outputs of which are summed in polarization pairs, high-pass filtered, integrated and one-bit digitized at intervals of 250 $\\mu$s. The observation time per pointing is 35 min. Using clusters of workstations at the various collaborating institutions, data are dedispersed with up to 325 trial dispersion measures (DMs) and searched for periodic signals. For low-DM pulsars with periods in the range 10 ms to 5 s, the limiting sensitivity of the survey is about 0.2 mJy. The survey commenced in mid-1997 and approximately 95\\% of the 2670 pointings required to complete it have been observed. Most of the data have been processed, resulting in the discovery of 608 pulsars, including eight binary pulsars and four millisecond pulsars. The survey and discovery of the first 100 pulsars are described in some detail by Manchester et al. (2001). Other papers announcing discoveries of special interest are referenced in that paper. ", + "conclusions": "" + }, + "0112/astro-ph0112216_arXiv.txt": { + "abstract": "We present an analysis of the X-ray emission of the supernova remnant \\subj, which was partially covered by three observations with \\xmm. The detection of Fe K emission at 6.4~keV, and the lack of spatial correlation between hard X-ray and radio emission is evidence against a dominant X-ray synchrotron component. We argue that the hard X-ray continuum is best explained by non-thermal bremsstrahlung from a supra-thermal tail to an otherwise cool electron gas. The existence of low electron temperatures is supported by low temperatures found in other parts of the remnant, which are as low as $0.2$~keV in some regions. ", + "introduction": "The X-ray emission from the supernova remnant \\subj\\ (RCW 86, G315.2-2.3) is characterized by spatially distinct soft and hard X-ray components (Vink, Kaastra, \\& Bleeker 1997). The soft X-ray emission has a thermal nature, but the hard X-ray emission shows relatively little line emission, suggesting X-ray synchrotron radiation (Borkowski et al. 2001), similar to SN 1006 (Koyama et al. 1995). Borkowski et al. argue that additional evidence for X-ray synchrotron radiation is the spatial correlation between hard X-ray and radio emission. A problem is, however, that the hard X-ray emission is accompanied by Fe K line emission at 6.4~keV (Vink et al. 1997). The energy of the line emission indicates that the emission is from underionized iron (Fe XVII or lower), but it also proofs that the line emitting plasma contains electrons with energies in excess of 7.1~keV. These electrons should give rise to bremsstrahlung emission, suggesting that at least part of the hard X-ray continuum is bremsstrahlung. We present here \\xmm\\ data of \\subj. \\xmm\\ offers a superior sensitivity and spatial resolution (FWHM $\\sim$ 6\\arcsec) compared to ASCA. This allows for a better separation of the hard and soft X-ray emitting regions. The complete set of observations covers most of the remnant, but here we only discuss data from the southeastern, central and northwestern part, which were observed in August 2000 with exposures of $\\sim 12$~ks to $17$~ks. More recent observations of the southwestern part will be presented in a future article. \\begin{figure} \\psfig{figure=vinkj1_1.ps,width=\\textwidth} \\caption{Soft X-ray (0.5-1 keV), left, and hard X-ray (2-7~keV) \\xmm\\ mosaics of \\subj, combining exposure and effective area corrected data from all three CCD cameras; radio contours (ATCA, Dickel, Strom, \\& Milne, 2000) are overlayed. \\label{mosaics}} \\end{figure} ", + "conclusions": "Based on \\xmm\\ observations we have argued that the hard X-ray emission from \\subj\\ is (non-thermal) bremsstrahlung rather than synchrotron emission. The mysterious X-ray emission may play a key role in studying non-thermal processes associated with collisionless shock heating and acceleration. To understand the nature of the hard X-ray emission, detailed maps of the continuum and Fe K emission are needed. A close correlation between the two would indicate a bremsstrahlung nature. A lack of correlation would indicate that the hard X-ray emission comes from two distinct components: ultrarelativistic electrons and a hot plasma. This issue can therefore be resolved by a deep \\xmm\\ observation." + }, + "0112/astro-ph0112020_arXiv.txt": { + "abstract": "Strong toroidal magnetic fields generated in stellar collapse can generate magneto-centrifugal jets in analogy to those found in simulations of black hole accretion. Magneto-centrifugal jets may explain why all core collapse supernovae are found to be substantially asymmetric and predominantly bi-polar. We describe two phases: the initial LeBlanc-Wilson jet and a subsequent protopulsar or toroidal jet that propagates at about the core escape velocity. The prompt LeBlanc-Wilson jets will produce an excess of neutron-rich matter and hence cannot be the common origin of supernova explosions; similar, but less severe problems arise with the protopulsar jet that may be alleviated by partial evacuation along the axis by rotation. The jets will produce bow shocks that tend to expel matter, including iron and silicon, into equatorial tori. This may help to account for observations of the element distribution in Cas A. A magnetic ``switch'' mechanism may apply in rare instances (low density and large magnetic field), with subsequent increase in the speed and collimation of the toroidal jet. The conditions that turn the magnetic switch ``on'' would yield a jet that propagates rapidly and with small opening angle through the star, depositing relatively little momentum. The result could be enough infall to form a black hole. A third, highly relativistic jet from the rotating black hole could catch up to the protopulsar jet after it has emerged from the star. The interaction of these two jets plausibly could be the origin of the internal shocks thought to produce \\grbs\\ and could explain the presence of iron lines in the afterglow. Recent estimates that typical \\grb\\ energy is $\\sim 3\\times10^{50}$ erg imply either a very low efficiency for conversion of rotation into jets by the Blandford-Znajek mechanism, or a rather rapid turnoff of the jet process even though the black hole still rotates rapidly. Magnetars and ``hypernovae'' might arise in an intermediate parameter regime of energetic jets that yield larger magnetic fields and provide more energy than the routine case, but that are not so tightly collimated that they yield failed supernova. ", + "introduction": "The problem of core-collapse has been with us for over 40 years (Hoyle \\& Fowler 1960). Immediately after the discovery of pulsars, it was reasonable to explore the issue of whether or not the rotation and magnetic fields associated with pulsars could be a significant factor in the explosion mechanism (Ostriker \\& Gunn 1971; Bisnovatyi-Kogan 1971; Bisnovatyi-Kogan \\& Ruzmaikin 1976; Kundt 1976). With typical dipole fields of $10^{12}$ G and rotation periods of several to several tens of milliseconds implying electrodynamic power of only $\\sim 10^{44-45}$ erg s$^{-1}$, a strong robust explosion seemed unlikely. Several factors have led to a need to re-examine this conclusion. The principle one is the accumulating evidence that core collapse supernovae are distinctly asymmetric. Aside from its famous rings, HST observations of SN~1987A resolving the debris show that the ejecta are asymmetric with an axis that roughly aligns with the small axis of the rings (Pun et al. 2001, Wang et al. 2002). Chandra X-ray Observatory (CXO) observations of Cas A show that the jet and counter jet and associated structure are observable in the X-ray (Hughes et al. 2000; Hwang et al. 2000) as well as the optical (Fesen \\& Gunderson 1996; and references therein). The most direct evidence bearing on this topic is from supernova spectropolarimetry which shows that substantial asymmetry is ubiquitous in core-collapse supernovae, and that a significant portion show strong evidence for a single, wavelength-independent axis of symmetry (Wang et al. 1996; Wang et al. 2001). Many, even perhaps most, core-collapse supernovae are bi-polar (Wang et al. 2001, 2002). The strength of the asymmetry observed with polarimetry is higher (several \\%) in supernovae of Type Ib and Ic that represent exploding bare non-degenerate cores (Wang et al. 2001). The degree of asymmetry also rises as a function of time for Type II supernovae that have retained their hydrogen envelopes (from $\\lta$ 1\\% to $\\gta$ 1\\%) as the ejecta expand and one looks more deeply into the core material (Wang et al. 2001a; Leonard et al. 2000, 2001). Both of these trends suggest that it is the inner machine, the core collapse mechanism itself, that is responsible for the asymmetry. The observed polarization requires significant asymmetry; axis ratios exceeding 2 to 1. To impose the observed strong asymmetry in the final homologous expansion it is plausible that an axial flow must be established and maintained for at least several dynamical time scales of the outer mantle/envelope. This is the operational definition of a jet. It also is significant that CXO observations have determined that jets are routinely associated with pulsars, not only in binary systems like SS 433, but also isolated objects like the Crab and Vela pulsars (Weisskopf et al. 2000; Helfand, Gotthelf, \\& Halpern 2001). One interpretation of these observations is that the present-day jets in these young objects are vestiges of a much more powerful MHD jet era that occurred in the first few seconds of the protopulsar phase when the compact objects were still inside their progenitor stars as they first attain nuclear densities. During this period the transient values of the magnetic field and rotation could have greatly exceeded those observed today, and indeed, as we suggest here, {\\em the reduction of the rotation rate and magnetic field to their present-day values could have been the process of energy release that powered the initial explosion}. Asymmetries associated with neutrino emission may produce dynamical asymmetries, but it is not clear that they can account for the polarization observations. Neutrino asymmetries may yield a short-lived, essentially impulsive effect (Shimizu, et al. 1994; Burrows \\& Hayes 1996; Fryer \\& Heger 2000; Lai et al. 2001). Expansion and transverse pressure gradients can wipe out transient asymmetries before homologous expansion is achieved (Chevalier \\& Soker 1989). ``Finger'' asymmetries might be preserved, but it is unclear that they can reproduce the common feature of a single symmetry axis that is substantially independent of space and time (Wang et al. 2001, 2002). Sufficient neutrino impulse might be delivered to the neutron star to yield a substantial runaway velocity (Burrows \\& Hayes 1996; Spruit \\& Phinney 1998), but it is difficult to see how this impulse can be communicated in a substantial and permanent way to the final ejecta trajectories. Highly-resolved, fully three dimensional, adaptive grid numerical calculations (Khokhlov 1998) have, however, established that non-relativistic axial jets of energy of order $10^{51}$ ergs originating within the collapsed core can initiate a bi-polar asymmetric supernova explosion that is consistent with the spectropolarimetry (Khokhlov et al. 1999; Khokhlov \\& H\\\"oflich 2001; H\\\"oflich, Khokhlov \\& Wang 2001). Some imbalance in axial jets can also account for pulsar runaway velocities, specifically velocities that are parallel to the spin axis (Helfand, Gotthelf \\& Halpern 2001, and references therein). While a combination of neutrino-induced and jet-induced explosion may prove necessary for complete understanding of core-collapse explosions, jets as computed by Khokhlov et al. are sufficient. In this paper, we will ignore the possibility of neutrino-induced explosions as we focus on the possibilities of jet formation and jet-induced supernovae. For purposes of discussion, we will consider the sort of structure that forms a stalled shock after core bounce. Further work on jet-induced supernovae has been done by Khokhlov \\& H\\\"oflich (2001) and H\\\"oflich, Khokhlov \\& Wang (2001). These papers present well-resolved, three-dimensional jets that propagate from within the region of the original iron core until breakout from the star or stoppage in a red giant envelope in a single calculation that does not require artificial halting, rescaling, and restarting the jet. A bow shock forms at the head of the jet and spreads roughly cylindrically around each jet. The stellar matter is shocked by the bow shock and acts as a high-pressure confining medium by forming a cocoon around the jet. The jets become long bullets of high-density material moving through the background low-density material almost ballistically. Spreading is limited by a secondary shock that forms around each jet between the jet and the material already shocked by the bow shock. The laterally expanding bow shocks generated by the jets move towards the equator where they collide with each other. The result is that the material in the equatorial plane is compressed and accelerated more than material in other directions (excluding the jet material). The result is that heavy elements (e.g. O, Ca) are characteristically ejected in tori along the equator. Iron, silicon and other heavy elements in Cas A are distributed in this way (Hwang et al. 2000), and there is some evidence for this distribution in SN 1987A (Wang et al. 2001b). The presence of a red giant envelope can stop the jet and dissipate the effects of the asymmetry in the outer envelope, but the asymmetry will still be manifest in the core. Radioactive matter ejected in the jets can alter the ionization structure and hence the shape of the photosphere of the envelope even if the density structure is spherically symmetric (H\\\"oflich, Khokhlov \\& and Wang 2001). This will generate a finite polarization, even though the density distribution is spherical and the jets are stopped deep within the star and may account for the early polarization observed in Type II supernovae (Leonard et al. 2000; Wang et al. 2001). The jets are unstable to Kelvin-Helmholtz instabilities, but the growth time is long compared to the jet propagation time and no significant modulation of the jet is observed. Faster jets tend to make weaker explosions by propagating so rapidly and narrowly through the star that there is relatively little time to generate the transverse shocks that cause the explosion. The opening half angle of the jet is approximately (Wheeler et al. 2000), \\begin{equation} \\label{jetangle} \\theta \\simeq \\frac{v_{env}}{v_{bow}}\\simeq 0.1~{\\rm rad} \\frac{v_{env,8}}{v_{bow,9}}\\simeq 5^o\\frac{v_{env,8}}{v_{bow,9}}. \\end{equation} where $v_{env}$ is the speed of sound of the envelope and $v_{bow}$ is the speed of the bow shock, which is less than that of the inflow velocity of the jet material. As $v_{bow}$ approaches the speed of light, the jet will be very narrow and affect a very small volume of the star as it propagates out. This feature may be relevant to the outcome of the explosion process and to the possibility of making \\grbs, as we will discuss in \\S 4.3. Radiative transfer effects within optically thin portions of the jet may lead to significant instabilities in faster, lower density, jets (H\\\"oflich, private communication). For related work in the context of black hole formation, see MacFadyen \\& Woosley 1999; Aloy, et al. 2000; MacFadyen, Woosley \\& Heger 2000; Zhang \\& Woosley 2001). The question then arises as to the mechanism of the production of the jets in routine core collapse events. A significant role for rotation and magnetic fields is an obvious candidate. Any mechanism that purports to account for routine pulsar formation may involve large transient magnetic fields, but must ultimately be consistent with the distribution of deduced dipole strengths of ``normal'' pulsars of $\\sim 10^{12}$--$10^{13}$ G and with the distribution of initial rotation periods. There is growing evidence for a considerable spread in these quantities ({\\it e. g.} Kaspi, et al. 2001 and references therein). In a minority of cases, the final dipole field might be consistent with a value of $\\sim 10^{15}$ G, yielding a ``magnetar'' (Duncan \\& Thompson 1992). Growing evidence for neutron stars of that field strength has been obtained by RXTE and other facilities (Kouveliotou et al. 1999; Ibrahim et al. 2001). The nature of the birth event of a magnetar is a separate problem that is significant in its own right. The ensemble of normal pulsar and magnetar births must also be consistent with observed nucleosynthetic abundances, a potentially crucial constraint on jet-induced supernova models that we discuss in \\S 4. Possible physical mechanisms for inducing axial jets, asymmetric supernovae, and related phenomena driven by magneto-rotational effects were considered by Wheeler et al. (2000). The means of amplifying magnetic fields by linear wrapping associated with differential rotation in the neutron star and possibly by $\\alpha$~--~$\\Omega$ dynamos was discussed. Attention was focused on the effect of the resulting net dipole field on the torquing of the infalling plasma and newly formed neutron star and on the creation of strong Poynting flux (in the form, initially, of ultrarelativistic MHD waves) at the time-variable speed of light circle in analogy to pulsar radiation mechanisms, albeit buried deeply in the core-collapse supernova ambience. Allusion was made to the existence and possible role of the primary toroidal field that is expected to form in routine collapse where protoneutron star rotation rates are not expected to trigger $\\alpha$~--~$\\Omega$ dynamos and exponential field growth. In this paper, we will explore the capacity of the predominantly toroidal field generated near the core-mantle boundary to directly generate axial jets within core-collapse conditions by analogy with magneto-centrifugal models of jets in AGN (Koide, Shibata \\& Kudoh 1997; Meier et al. 1997; Romanova et al. 1998; Ustyugova et al. 1999; Meier, 1999; Koide, Meier, Shibata \\& Kudoh 2000; and references therein). The magneto-centrifugal models in the literature usually focus on black hole conditions where the field is anchored in the disk and at infinity. The physical picture discussed in this paper may prove especially robust since the field is anchored in substantial, corporeal objects: the outer, still collapsing layers of the progenitor star and the collapsed object --- a neutron star. We find that the production of a strong toroidal field, substantially stronger than the $10^{12}$ G dipole field of a pulsar, is nearly inevitable, and strong axial jets driven by that field equally so. The mechanisms described here also may prove to have a close astrophysical analog in the production of supersonic magneto-centrifugal jets in young stellar objects where also there is an outer envelope of infalling matter and an inner, differentially rotating object (Konigl \\& Pudritz 2000; Lery \\& Frank 2000; Ouyed, Pudritz \\& Stone 1997). Our working hypothesis is thus that jets leading to asymmetric supernova explosions are associated with routine core collapse and pulsar formation in roughly 90\\% of the cases. Perhaps 10\\% of the core collapse events would be associated with magnetars. Magnetar formation might be associated with exceptionally strong explosions and asymmetries, but this is not necessarily the case. An even smaller fraction, perhaps $10^{-3}$ to $10^{-4}$ (Scalo \\& Wheeler 2002a) might involve black hole formation and the production of \\grbs. In \\S 2 we outline the circumstances by which a newly formed neutron star could yield MHD jets. In \\S 3, we discuss the basic physics of MHD jets and their formation. Section 4 explores the application of the physical principles of jet formation to supernova conditions, how magneto-centrifugal jets could be generated in core collapse, and some of their expected properties. We sketch a scenario in which extreme values of the rotation and magnetic field could lead to the failure of the initial supernova explosion with the subsequent collapse of the core to form a black hole and the attendant possibility to form a \\grb. In \\S 5 we summarize our conclusions. ", + "conclusions": "We have presented the case that strong toroidal magnetic fields will be generated in stellar collapse and that these magnetic fields can generate magneto-centrifugal jets in analogy to those found in simulations of black hole accretion. The case for magneto-centrifugal jets suggests they may be frequent and robust and hence provide a good basic understanding of why all core collapse supernovae are found to be substantially asymmetric and predominantly bi-polar. There are concerns that the jets we describe arise deep in the gravitational potential where the infalling matter is very neutron rich. We argue that this constraint rules out the prompt LeBlanc-Wilson jets as the common origin of supernovae. This suggests that conditions of initial rotation and magnetic field extreme enough to produce strong jets in this phase do not arise frequently in nature. We point out that successful jets from the later, protopulsar phase must walk the line between being broad enough to provide ample energy to the overlying matter without being so broad as to eject excessive neutron-rich matter. Partial evacuation of the matter on the axis by rotation may alleviate this problem, but it remains a concern. The jets will produce bow shocks that tend to expel matter, including iron and silicon, in equatorial tori. This may help to account for observations of the element distribution in Cas A. There also is some concern that gravitational radiation may remove so much angular momentum, and so quickly, that there would not be enough rotational energy to power a substantial MHD jet. We estimate that the effects of gravitational radiation (GR) will be significant, but not severe. The most important GR mode is probably the $r$-mode. Detailed computations by Lindblom, Tohline, and Vallisneri (2001) on the evolution and effects of this mode estimate that only about half of the rotational kinetic energy will be lost before the gravitational radiation process switches off. The time scale for this to occur may be fairly short (a few tens of initial rotational periods, i.e. as short as, say, 30 ms), but the protoneutron star still will be left with a significant fraction of its initial rotational energy, which can be tapped for the production of MHD jets. The topics of core-collapse supernovae and \\grbs\\ overlapped, at least in principle, with the discovery of SN 1998bw. Although the association is still controversial, this odd supernova is likely to have been connected to GRB~980425 (Galama et al. 1998). Supernova-like excesses of light have been detected in the afterglow of two \\grbs, GRB~970228 (Reichart 1999, Galama et al. 2000) and GRB~980326 (Bloom et al. 1999) about two weeks after the \\grbs. The excess light in the afterglow of both GRB~970228 and GRB~980326 has been modeled by the addition of light from an event like SN~1998bw. While there may be other explanations for these excesses, the connection with supernovae must be pursued. We suggest here a new way to make this connection. The conditions that apply in stellar collapse suggest that a magnetic ``switch'' mechanism found to apply in the black hole simulations may also apply in the collapse case with subsequent affect on the speed of the jet. If the switch turns on at low density and large magnetic field, an especially fast jet, with propagation speed of order the Alfv\\'en speed, could be produced. This jet would tend to propagate rapidly through the star and deposit relatively little momentum. The result might be a supernova explosion, but with enough infall to produce a delayed black hole. A reprise of the physics outlined here, and that explicitly invoked to produce relativistic MHD jets from stellar mass and AGN black holes, could then produce a relativistic jet that catches up to the first jet after it has emerged from the star. The interaction of these two jets could plausibly be the origin of the internal shocks thought to produce \\grbs\\ and could explain the presence of iron lines in the afterglow. In this paper we have discussed toroidal fields of order $10^{14}$ to $10^{16}$ G arising routinely in the formation of a neutron star in a supernova. There is clearly an issue of the final effective dipole strength of the neutron stars left behind in the explosion. We must have a final dipole field distribution of the neutron stars consistent with those of observed pulsars, $\\sim$ $10^{12}$ to $10^{13}$ G. It seems plausible that the large fields generated in the collapse are dissipated in the helical jets. The very strong fields that may be generated in the depths of the neutron star will take longer to float out, but the time still could be short compared to the lifetime of the supernovae, never mind the pulsar. Stronger fields in the interior could also leave smaller surface dipole fields. Much of the field may be amplified and dissipated at the shearing boundary between the infalling matter and the neutron star. This layer and its field will almost surely substantially dissipate after the explosion. These issues should be given more careful consideration. All of the issues raised here should be examined in the context of magnetars. However the large effective dipole fields arise in magnetars, there is a strong presumption that they are created in the formation of the neutron star. The formation of a magnetar may require special circumstances, or they may arise as the tail of the normal pulsar formation process. One possibility is that they represent that fraction of core-collapse events that, perhaps by accident of progenitor conditions, do attain an especially rapid differential rotation and amplify the field not just by field line wrapping, but by an $\\alpha - \\Omega$ dynamo (Duncan \\& Thompson 1992). This could lead to exponential field growth, and total field strength of order $10^{17}$ G with dipole component of order $10^{15}$ G. This is approximately the regime of fast toroidal jets that we invoke to produce less ejecta and more infall to produce a black hole and perhaps a \\grb. To reconcile these pictures, it may be that magnetars arise with somewhat higher than normal fields and rotation, but still in the slow toroidal jet phase with conditions where the magnetic switch is still ``off.\" Even more extreme conditions could lead to the flipping of the switch, producing an especially fast jet that would be less effective in driving the explosion and thus leaving a black hole. Clearly there is a great need to study the many aspects of this problem in numerical and quantitative detail. Some models of SN~1998bw invoked especially large kinetic energies, in excess of $10^{52}$ ergs in spherically-symmetric models, to account for the bright light curve and high velocities (Iwamoto, et al. 1998; Woosley, Eastman, \\& Schmidt 1998), while others took note of the measured polarization to suggest that strongly asymmetric models could account for the observations with more ``normal'' energies (H\\\"oflich, Wheeler, \\& Wang 1999). Large kinetic energy has also been attributed to several other supernovae (again based on spherically symmetric models), especially SN 1997cy (Germany et al. 2000) and SN 1997ef (Branch 2001; Iwamoto et al. 2000). Events like SN~1997cy, SN~1997ef and SN~1998bw will help to sort out the physics of explosive events, whether such events are more closely related to ``ordinary'' supernovae or ``hypernovae'', whether either of these classes leaves behind neutron stars as ``ordinary'' pulsars or highly magnetized ``magnetars'' or whether the remnant is a black hole and whether any of these events are associated with classic cosmic \\grbs\\ as suggested by the supernova-like brightening of the afterglow of GRB~970228 and GRB~980326. In terms of our magnetic switch mechanism, it may be that ``hypernovae\" are still in the switch ``off\" phase so they have strong jets, but not so fast and narrow that the explosion is mitigated. In this sense, there may be a connection between the ``hypernovae\" and magnetars that, as argued above, may arise in the same regime of especially energetic protopulsar jets, but still with the switch ``off.\" If it were not for the constraints of the supernova polarimetry, we might continue to study essentially spherically symmetric collapse models, as indeed, many people will and must. It may be that we are pessimistic and that neutrino asymmetries can account for the observed polarimetry. The fact that jets alone can account for supernova explosions and the observed asymmetries has been established in principle. We think nature is saying that jets must occur in routine supernovae. We also believe that the process of neutron star formation must inevitably involve the physics outlined here whether the resulting jets are weak or strong. Clearly, we need a more rigorous description of the process of field generation, buoyancy, emergence, and subsequent evolution." + }, + "0112/astro-ph0112034_arXiv.txt": { + "abstract": "We present a study of $HST-WFPC2$ observations of the inner kpc of the interacting galaxy M51 in six bands from 2550 \\AA\\ to 8140 \\AA. The images show an oval shaped area (which we call ``bulge\") of about $11 \\times 16$ arcsec or $450 \\times 650$ pc around the nucleus that is dominated by a smooth ``yellow/reddish\" background population with overimposed dust lanes. These dust lanes are the inner extensions of the spiral arms. The extinction properties, derived in four fields in and outside dust lanes, is similar to the Galactic extinction law. The reddish stellar population has an intrinsic color of $(B-V)_0 \\simeq 1.0 $ suggesting an age in excess of 5 Gyrs. We found 30 bright point-like sources in the bulge of of M51 i.e. within 110 to 350 pc from the nucleus. The point sources have $21.4 < V < 24.3$, many of which are blue with $ B-V < 0$ and are bright in the UV with $19.8 < m_{2550}< 22.0$. These objects appear to be located in elongated ``strings\" which follow the general pattern of the dust lanes around the nucleus. The spectral energy distributions of the point-like sources are compared with those predicted for models of clusters or single stars. There are three reasons to conclude that most of these point sources are isolated massive stars (or very small groups of a few isolated massive stars) rather than clusters:\\\\ (a) The energy distributions of most objects are best fitted with models of single stars of $M_V$ between -6.1 and -9.1, temperatures between 4000 and 50000 K, and with 4.2 $<$ log $L/\\Lsun <$ 7.2, and $12 < \\Mstar < 200~ \\Msun$. \\\\ (b) In the HR diagram the sources follow the Humphreys-Davidson luminosity upper limit for massive stars. \\\\ (c) The distribution of the sources in the HR diagram shows a gap in the range of $20~000 < \\Teff < 10~000$ K, which agrees with the rapid crossing of the HRD by stars, but not of clusters.\\\\ We have derived upper limits to the total mass of lower mass stars ($\\Mstar < 10~ \\Msun$), that could be ``hiding'' within the point sources. For the ``bluest'' sources the upper limit is only a few hundred \\Msun. We conclude that the formation of massive stars outside clusters (or in very low mass clusters) is occurring in the bulge of M51. The estimated star formation rate in the bulge of M51 is 1 to $2 \\times 10^{-3}$ \\Msunyr, depending on the adopted IMF. With the observed total amount of gas in the bulge, $\\sim 4 \\times 10^5$ \\Msun, and the observed normal gas to dust ratio of $\\sim 150$, this star formation rate could be sustained for about 2 to $4 \\times 10^8$ years. This suggests that the ongoing massive star formation in the bulge of M51 is fed/triggered by the interaction with its companion about $4 \\times 10^8$ years ago. The star formation in the bulge of M51 ia compared with tat in bulges of other spirals. Theoretical predictions of star formation suggest that isolated massive stars might be formed in clouds in which H$_2$, [OI] 63 $\\mu$m and [CII] 158 $\\mu$m are the dominant coolants. This is expected to occur in regions of rather low optical depth, $A_V \\le 1$, with a hot source that can dissociate the CO molecules. These conditions are met in the bulge of M51, where the extinction is low and where CO can be destroyed by the radiation from the bright nuclear starburst cluster in the center. The mode of formation of massive stars in the bulge of M51 may resemble the star formation in the early Universe, when the CO and dust contents were low due to the low metallicity. ", + "introduction": "The spiral galaxy M51 (NGC 5194, the Whirlpool nebula) and its peculiar companion galaxy NGC 5195 form a typical example of galaxy interactions. After the pioneering hydrodynamical simulations by Toomre and Toomre (1972), several authors have tried to explain the grand design spiral shape and the tidal arms of this interacting system. Hernquist (1990) and Barnes (1998) have critically discussed the successes and the problems of explaining the morphology of the NGC 5194/5195 system, in particular the radial velocities of the two galaxies, the two-armed spiral structure of M51, the connecting tidal arm and the large H\\,{\\sc i} arm. The best model is found for a relative orbit that is almost in the plane of M51, for a mass ratio of NGC 5194/5195 $\\simeq 2$, a pericenter distance of 17 to 20 kpc and a time since pericenter of $2.5\\times 10^8 $ 5 Gyrs) stars (Paper I). The full $HST-WFPC2$ image of M51 is published is Paper I. The image shows that the nucleus is surrounded by an elongated bulge of about 460 $\\times$ 860 pc that is dominated by an old stellar population. The spiral arms containing H\\,II regions start outside the bulge. We adopt a distance of $d=8.4 \\pm 0.6$ kpc, corresponding to a distance modulus of $29.62 \\pm 0.16$, which is based on the brightness distribution of planetary nebulae (Feldmeier, Ciardullo \\& Jacoby 1997). At this distance, 1 arcsec corresponds to a linear distance of 40.7 pc, one $HST-PC$ pixel of 0.046$^\"$ corresponds to 1.87 pc and an $HST-WFC$ pixel of 0.1$^\"$ corresponds to 4.1 pc. In \\S~2 we describe the observations and the data reduction. In \\S~3 we discuss the interstellar extinction in the bulge. In \\S~4 we describe the properties of 30 bright point-like sources in this region. In \\S~5 their energy distributions are compared with those predicted for clusters with different ages and mass, and for single stars of different effective temperatures and radii. We will show that most of them are very luminous young stars (single or multiple), rather than clusters. The star formation rate in the bulge of M51 is derived in \\S~6. In \\S~7 we compare the formation of massive stars in the bulge of M51 with clusters near the Galactic center, in the interaction region of the Antennae galaxies, and in the bulges of other spiral galaxies. We also discuss the predicted mode of star formation under the conditions that prevail near the nucleus of M51. The conclusions are given in \\S~8. ", + "conclusions": "We have studied bright point sources in the bulge of M51 with the HST-WFPC2 camera in 6 filters in the wavelength region of 2500 to 8200 \\AA. The results can be summarized. \\begin{enumerate} \\item We found 30 point sources in the bulge of M51 with $21.38 < V < 26.20$. These point sources appear to occur in strings that follow the general pattern of the elliptical or spiral-like dust lanes in the bulge of M51, but they do not necessarily coincide with the dust lanes. \\item The extinction of the point sources, derived by fitting their energy distribution with those of single star models or cluster models, is in the range of $0 < E(B-V) < 1.3$. Half of the objects have $E(B-V)<0.25$. The absolute visual magnitudes range from about $M_V \\simeq-6$ to $-9$. \\item The energy distributions and the distribution of the point sources in the HR diagram suggest that most of them, except one or two, are stars rather than clusters because:\\\\ (a) The observed energy distributions better fit those of stellar models than those of cluster models.\\\\ (b) Many of the sources have an energy distribution and a luminosity that is characteristic for a single hot massive star of $25~000 \\le \\Teff \\le 50~000$ K. \\\\ (b) The distribution of the objects in the HR diagram follows the Humphreys-Davidson luminosity upper limit for stars. There is no reason why clusters would follow this limit. \\\\ (c) There is a gap in the distribution of the sources in the HR diagram at intermediate temperatures between $\\Teff \\simeq 20~000$ and 10~000 K. This is easily explained if the point sources are stars, because that temperature range is crossed rapidly by the evolution tracks of stars. There is no obvious reason why clusters would avoid this temperature or colour range. \\item The distribution of stars in the HR diagram shows two groups: (a) the most massive group of initial mass $M_i > 40~ \\Msun$ is mainly blue and hot, because most of these stars will not evolve into red supergiants during their evolution. (b) the group with $12 < M_i < 25~ \\Msun$ is mainly red and cool because stars in this mass range are below the detection limit during their main sequence phase. \\item The current star formation rate in the bulge of M51 in the mass range of $12 < M < 200$ \\Msun\\ is $\\sim 5 \\times 10^{-4}$ \\Msunyr. Correcting for the possible presence of lower mass stars, down to 1 \\Msun, increases the star formation rate by a factor 3.4 if we adopt an initial mass function of slope $\\Gamma = -1.35$ (Salpeter's value) and by a factor 1.5 if we adopt $\\Gamma=-0.65$ (the value for the clusters near the Galactic center). \\item The total amount of neutral H in the bulge is about $4 \\times 10^5$ \\Msun\\ and the gas-to-dust mass ratio about 150. The current star formation rate of about $2 \\times 10^{-3}$ \\Msunyr\\ can be sustained for about 2 to $4 \\times 10^8$ years before all the gas is consumed. This suggests that this form of massive star formation in the bulge is fed/triggered by the interaction with the companion galaxy, whose closest approach was estimated to be about $4 \\times 10^8$ years ago. \\item These results show that under the conditions that exist in the bulge of M51, separate massive stars can form outside clusters or in very small groups. This agrees with the predictions of Norman and Spaan (1997) who argued that the formation of massive stars is favoured in regions near a hot source (the core of M51) and small optical depth of $A_V \\le 1$ so that CO is dissociated but H$_2$ survives due to self-shielding. This may resemble the star formation in the early Universe, when the CO content and the dust content were low due to the low metallicity. \\end{enumerate}" + }, + "0112/astro-ph0112172_arXiv.txt": { + "abstract": "The ANTARES collaboration is building a deep sea neutrino telescope in the Mediterranean Sea. This detector will cover a sensitive area of typically 0.1~km$^2$ and will be equipped with about 1000 optical modules. Each of these optical modules consists of a large area photomultiplier and its associated electronics housed in a pressure resistant glass sphere. The design of the ANTARES optical module, which is a key element of the detector, has been finalized following extensive R\\&D studies and is reviewed here in detail. ", + "introduction": "\\label{intro} Neutrinos offer a unique opportunity to explore the Universe in depth over a wide energy range~\\cite{bib:general}. However, since neutrino fluxes at high energy (E$_\\nu>$~TeV) are expected to be very low~\\cite{bib:theoflux}, a very large detector volume is required. A detector immersed in the sea provides a cheap and efficient method of observing high energy muon neutrinos by detecting the muon produced from a charged current interaction in the matter surrounding the detector. The muon emits \\v{C}erenkov light as it passes through the water, and this light can be detected by a three-dimensional array of optical sensors called Optical Modules (OM). The measurement of the arrival time of the \\v{C}erenkov light at each OM allows the reconstruction of the muon direction, and the amount of light collected can be used to estimate the muon energy. The ANTARES collaboration has started the construction of a \\proto\\ detector to be immersed at 2400~m depth in the \\med\\ 40~km off the French coast (42$^0$ 50' N, 6$^0$ 10' E). This detector will be equipped with about 1000 OMs and has been designed to be competitive with other such detectors (DUMAND~\\cite{bib:dumand}, Baikal~\\cite{bib:baikal}, AMANDA~\\cite{bib:amanda} and NESTOR~\\cite{bib:nestor}), notably in terms of angular resolution. To reach this goal, extensive R\\&D studies have been carried out during the past years in order to optimize the design of each element of the telescope, particularly its most fundamental component, the optical module. This optimization work has involved a variety of measurements, both in the laboratory and {\\it in situ}, as well as Monte Carlo simulations. This paper gives the results of these studies and describes in detail the final design of the ANTARES optical module. In section~\\ref{sec:omspecif}, the constraints inherent to the operation of a deep sea neutrino detector are reviewed and the specifications for the optical modules are listed. Section~\\ref{sec:omassembly} describes the main features of the different components of the OM as well as the assembly procedure. The test set-ups and the main results of the tests performed by the collaboration during the past years are presented in section~\\ref{sec:omtest}. Finally, section~\\ref{sec:massprod} deals with mass production and quality control aspects. ", + "conclusions": "The optical module is a key component of the ANTARES detector. In view of its importance, extensive R\\&D studies have been carried out by the collaboration to ensure that it would not only be optimized in terms of performance, but would also be robust enough to withstand sea-operation stresses and be highly reliable so as to guarantee proper operation during the entire lifetime of the detector. The general design and the choice of components have been finalized, so that mass production of the OMs can begin." + }, + "0112/astro-ph0112491_arXiv.txt": { + "abstract": "The equation of state for compact stars is reviewed with special emphasis on the role of strange hadrons, strange dibaryons and strange quark matter. Implications for the properties of compact stars are presented. The importance of neutron star data to constrain the properties of hypothetic particles and the possible existence of exotic phases in dense matter is outlined. We also discuss the growing interplay between astrophysics and heavy-ion physics. ", + "introduction": "In the last few years, it became clear that strangeness has to be included as another degree of freedom in astrophysical systems. In this review, we will outline some recent developments in the study of matter with strangeness under extreme conditions with relevance to astrophysics. The field is growing rapidly, theoretically as well as experimentally, so we will not give an overview of the field but rather focus on some recent works about the properties of compact stars with a remark about the cosmological phase transition including strangeness. The topics we are going to cover are: hyperons in neutron stars, kaon condensation in hadronic matter, H-dibaryon condensation, strange quark phase and twins, kaon condensation in quark matter, signals from proto-neutron stars with a strange phase, strange quark stars, new data from a isolated neutron star and the cosmological phase transition with strangeness. The phase diagram of Quantum Chromodynamics (QCD) can be studied at large temperature and zero (or small) density by lattice calculations and by relativistic heavy-ion collisions. Neutron stars on the other hand probe QCD at high density and small temperature and provide therefore a complementary laboratory to study QCD under extreme conditions. Created by supernova explosions, neutron stars are compact remnants with masses around 1--2 solar masses and radii of about $R\\sim 10$ km. The central density of the compact star will be then several times normal nuclear matter density $n_0$. More than 1000 pulsars, rotating neutron stars, are known today. The most precisely measured mass is the Hulse-Taylor pulsar with $M=(1.4411\\pm0.00035)M_\\odot$. ", + "conclusions": "" + }, + "0112/astro-ph0112458_arXiv.txt": { + "abstract": "The Planck High Frequency Instrument (HFI) is the most sensitive instrument currently being built for the measurement of Cosmic Microwave Background anisotropies. In addition to unprecendented sensitivity to CMB temperature fluctuations, the HFI has polarisation-sensitive detectors in 3 frequency channels (143, 217 and 353 GHz), which will constrain full-sky polarised emission of the CMB and foregrounds at these frequencies. The sensitivity of the instrument will allow a clear detection of CMB polarisation signals and should yield a precise measurement of its power spectrum at all angular scales between $\\ell = 50$ and $\\ell = 1000$, as well as constraints on the polarised emission at larger scales where a polarised signal from inflationary gravity waves or from reionisation is expected in many cosmological scenarios. ", + "introduction": "The COBE-DMR \\cite{dmr-sci1}, Boomerang \\cite{boomerang-sci1,boomerang-sci2}, DASI \\cite{dasi-sci1,dasi-sci2} and Maxima \\cite{maxima-sci1,maxima-sci2,maxima-sci3} experiments together now have yielded strong constraints on the Cosmic Microwave Background (CMB) anisotropy power spectrum \\cite{jaffe-2001}, and in particular a convincing detection of the first three acoustic peaks, providing compelling evidence that indeed the primordial fluctuations were produced during an inflationary phase. The next great challenge in the the study of the statistical properties of the CMB is to measure the polarisation signal, both that due to acoustic oscillations in causally connected regions of the Universe before decoupling and, even more challenging, the primordial polarisation spectrum due to gravity waves generated during inflation. Measurement of the correlated spectrum of polarisation and temperature at sub-degree scales will provide yet another test of the basic acoustic oscillation scenario and of the global paradigm, while measurement (or absence of detection) of polarisation signals due to tensor modes could strongly constrain inflationary models, as well as yield direct observational evidence for the existence of primordial gravity waves \\cite{zaldarriaga-1997,kamionkowski-1997a,seljak-1997,kamionkowski-1997b}. This has been widely recognised by the CMB community, and a large number of experiments dedicated to detecting the CMB polarisation are currently in operation or being planned. In this paper, we review the overall design of the Planck High Frequency Instrument (HFI) as a polarisation sensitive instrument, and discuss its capabilities in terms of measuring CMB polarisation. ", + "conclusions": "Originally planned and proposed solely as a CMB temperature anisotropy sensitive instrument, the Planck HFI has been revised to become as well a CMB polarisation sensitive instrument. Unprecedented sensitivity at high resolution is achieved through the combination of the use of new polarisation sensitive bolometers, cooled to 100 mK on a spaceborne mission with low background, and of the selection of observing frequency bands where diffraction limits the resolution at the 5 to 7 arcminute level. In addition, the selected frequency range is at the expected minimum of the polarised emission from galactic foregrounds and extragalactic compact sources. A substantial effort is being made to understand the impact of all possible systematic instrumental effects throughout the detection process and the data reduction pipeline. The minimisation of such systematic errors through a rigorous choice of the instrumental setup, careful on-ground testing, and continuing dedicated effort to the development and optimisation of data reduction methods for polarisation measurements with the Planck HFI, give us confidence that the instrument can meet its ambitious objectives." + }, + "0112/astro-ph0112344_arXiv.txt": { + "abstract": "SS\\,433 is a jet emitting X-ray binary surrounded by the W50 radio nebula. The SS\\,433\\,/\\,W50 system is an excellent laboratory for studying relativistic jet interaction with the surrounding interstellar medium. In this context, part of the W50 nebula has been mapped with ISOCAM at 15\\,$\\mu$m. I will show the results particularly on the W50 west lobe, and on 2 emitting zones detected with IRAS who have also been observed in millimetre wavelength (CO(1-0) transition), and for one of them by spectroscopy with ISOLWS and ISOSWS between 2 and 200 $\\mu$m. ", + "introduction": "SS\\,433 is an X-ray binary probably composed of a high-mass star and a neutron star, with a 13 days binary period. The system emits relativistic (0.26\\,c) jets showing a 162.5 days precession observed at subarcsecond scale in radio till \\mbox{$\\sim 10^{17}$\\,cm} from SS\\,433. At large scale these jets are observed in X-ray begining at $\\sim 15'$ (14\\,pc) from SS\\,433, and they are responsible for the unusual elongated shape of W50 ($\\sim 1^\\circ \\times 2^\\circ$), the possible supernova remnant radio nebula around SS\\,433. We mapped at 15\\,$\\mu$m with \\mbox{ISOCAM}\\footnote{The ISOCAM data presented in this paper was analysed using \"CIA\", a joint development by the ESA Astrophysics Division and the ISOCAM Consortium. The ISOCAM Consortium is led by the ISOCAM PI, C. Cesarsky.}, the infrared camera on board of the Infrared Space Observatory (ISO), a small part of the eastern lobe where an X-ray knot lies, the north-east quarter of the central circular part of W50, and nearly all the western lobe. ", + "conclusions": "" + }, + "0112/astro-ph0112564_arXiv.txt": { + "abstract": "Many self-gravitating systems often show scaling properties in their mass density, system size, velocities and so on. In order to clarify the origin of these scaling properties, we consider the stationary state of N-body system with inverse power law interaction. As a simple case, we consider the self-similar stationary solution in the collisionless Boltzmann equation with power law potential and investigate its stability in terms of a linear symplectic perturbation. The stable scaling solutions obtained are characterized by the power index of the potential and the virial ratio of the initial state. It is suggested in general that the nonextensive system has much various stable scaling solutions than those in the extensive system. ", + "introduction": "introduction} There are many self-gravitating systems which are characterized by some scaling properties. For example, the inter stellar medium shows that its velocity dispersion $\\sigma$ is power law related with the system size $L$ or mass $M$\\cite{larson81} ($\\sigma\\sim L^{0.38}\\sim M^{0.2}$) and isothermal contour is characterized by the fractal dimension $D\\sim1.36$\\cite{falgarone91}. The observations by the Hubble Space Telescope show elliptical galaxies have a power law density distribution $\\rho\\sim r^{-n}$ (at outer region, $n\\sim 4$ and at inner region, $n\\sim 0.5-1.0$ for the bright elliptical galaxies and $n\\sim 2$ for the faint ones\\cite{merritt96}.). The distribution of the galaxies and the cluster of galaxies can be characterized by the fractal dimension $D\\sim 2$\\cite{pietronero}. In cosmological simulations based on the standard cold dark matter scenario, the density profile shows a power law distribution (at outer region, $n\\sim 3$ and at inner region, $n\\sim 1.0-1.5$\\cite{NFW97,makino01}.). Recently, in order to study the statistical properties of a self-gravitating system, we proposed the self-gravitating ring model\\cite{sota}, where all particles are moving, on a circular ring located in three-dimensional space, with mutual interaction of gravity in three-dimensional space. The numerical simulation shows that the system at the intermediate energy scale, where the specific heat becomes negative, has some peculiar properties such as non-Gaussian and power law velocity distribution($f(v)\\sim v^{-2}$), scaling mass distribution, and self-similar recurrent motion. In this model, the {\\it halo} particles which belong to the intermediate energy scale are considered to play an important role in realizing such specific characters. We are interested in the origin of these scaling properties from statistical mechanical point of view. In order to study the statistical properties of long range interaction such as gravity, Ising model, and spin glass, the model with power law potential has been used and revealed anomalous properties\\cite{Ispolatov2001,Campa2001,Campa2002}. For example, a gravitational-like phase transition\\cite{Ispolatov2001}, reduction of mixing\\cite{Campa2001}, and long relaxation\\cite{Campa2002} are observed. Using a model with an attractive $1/r^\\alpha$ potential in general D-dimensional space, we can control the extensivity of the system and the specific heat by changing the spatial dimension $D$ and the exponent of inverse power of the potential $\\alpha$. In this paper, we study the quasi-equilibrium state of N-body system with a power law potential. As a first step, we consider the collisionless Boltzmann equation (CBE) in place of N-body system and derive the self-similar stationary solution of CBE which has a scaling property appearing in the quasi-equilibrium state and discuss the linear stability by the use of energy functional analysis\\cite{kandrup90,kandrup91,perez96,Rey96}. In section \\ref{sec:N}, we show some general properties of N-body systems with power law potential. In section \\ref{sec:SS}, we derive the self-similar stationary solution of CBE with an attractive $1/r^\\alpha$ potential assuming spherical symmetry and isotropic orbit case in D-dimensional space. Stability for the linear perturbation around the self-similar stationary solution is investigated in section \\ref{sec:STABLE}. Section \\ref{sec:DISS} is devoted to discussion. ", + "conclusions": "" + }, + "0112/astro-ph0112087_arXiv.txt": { + "abstract": "We propose to identify pulsar-wind bubbles (PWBs) as the environment in which the afterglow emission in at least some gamma-ray burst (GRB) sources originates. Such bubbles could naturally account for both the high fraction of the internal energy residing in relativistic electrons and positrons ($\\epsilon_e$) and the high magnetic-to-internal energy ratio ($\\epsilon_B$) that have been inferred in a number of sources from an interpretation of the afterglow emission as synchrotron radiation. GRBs might occur within PWBs under a number of scenarios: in particular, in the supranova model of GRB formation a prolonged (months to years) period of intense pulsar-type wind from the GRB progenitor precedes the burst. Focusing on this scenario, we construct a simple model of the early-time structure of a plerionic supernova remnant (SNR), guided by recent results on the Crab and Vela SNRs. The model is based on the assumption of an ``equipartition'' upper bound on the electromagnetic-to-thermal pressure ratio in the bubble and takes into account synchrotron-radiation cooling. We argue that the effective upstream hydrogen number density for a relativistic shock propagating into the bubble is given by $n_{\\rm H, equiv}= [4p + (B^\\prime+{\\mathcal{E}}^\\prime)^2/4\\pi]/ m_p c^2$, where $B^\\prime$ and ${\\mathcal{E}}^\\prime$ are, respectively, the comoving magnetic and electric fields and $p$ is the particle pressure. We show that, for plausible parameter values, $n_{\\rm H, equiv}$ spans the range inferred from spectral fits to GRB afterglows and that its radial profile varies within the bubble and may resemble a uniform interstellar medium or a stellar wind. We consider how the standard expressions for the characteristic synchrotron spectral quantities are modified when the afterglow-emitting shock propagates inside a PWB instead of in a uniform interstellar medium and demonstrate that the predictions for the empirically inferred values of $\\epsilon_e$ and $\\epsilon_B$ are compatible with the observations. Finally, we outline a self-consistent interpretation of the X-ray emission features detected in sources like GRB 991216 in the context of the supranova/PWB picture. ", + "introduction": "{1\\over r^2}{d\\over dr}(r^2 \\gamma\\beta w) - \\gamma\\beta{dp\\over dr} + {\\gamma({\\mathcal{E}}-\\beta B)\\over 4\\pi r}{d\\over dr}(rB)=-{\\Lambda\\over c} \\end{equation} (the entropy equation), where in both cases we took account of the constancy of the product $r{{\\mathcal{E}}}$ inside the bubble. If $\\sigma_w$ is not $\\ll 1$, then most of the bubble volume will be in the equipartition regime ($\\delta \\approx 1$, or, equivalently, $\\epsilon_B\\approx \\epsilon_e$), in which case $\\Lambda$ will typically be dominated by synchrotron radiation.\\footnote{In the equipartition region, synchrotron self-Compton emission cannot exceed the synchrotron radiation under any circumstances: it is comparable to the synchrotron emission if the bubble is highly radiative, but it remains much smaller if the radiative cooling time is longer than the bubble expansion time.} To simplify the treatment, we take synchrotron emission to be the main radiative cooling process even for low values of $\\sigma_w$. Furthermore, we assume that at any given location within the bubble the $e^\\pm$ pairs have a monoenergetic energy distribution characterized by a random (or ``thermal'') Lorentz factor $\\gamma_e$. The latter approximation is appropriate if the postshock gas undergoes significant radiative cooling (e.g., \\cite{GPS00}), which, as we discuss in \\S \\ref{results}, may be the case in SMNS-driven bubbles. The synchrotron emissivity can then be written in the form \\begin{equation}\\label{Lambda} \\Lambda={4\\over 3}\\sigma_T c n \\gamma_e^2 {{B^\\prime}^2\\over 8\\pi}\\ , \\end{equation} where $\\sigma_T$ is the Thomson cross section. In view of equation (\\ref{rel_eos}), it is then also possible to write the particle pressure as \\begin{equation}\\label{p_e} p={1 \\over 3} \\gamma_e n m_e c^2\\, . \\end{equation} Combining this expression with equation (\\ref{particle2}) gives \\begin{equation}\\label{gamma_e} \\gamma_e= Dr^2\\gamma\\beta p\\, , \\end{equation} where $D \\equiv 3/m_ec^2C=3/\\gamma_{\\rm ps}\\beta_{\\rm ps}n_{\\rm ps}R_s^2m_ec^2$. Using equations (\\ref{EB}), (\\ref{C}), (\\ref{p_e}), and (\\ref{gamma_e}) in equation (\\ref{Lambda}), one can express the radiative cooling term in equation (\\ref{entropy}) in the form \\begin{equation}\\label{Lambda2} {\\Lambda\\over c}= G \\gamma^3\\beta(rB-\\beta r{\\mathcal{E}})^2 p^2\\, , \\end{equation} where $G\\equiv (\\sigma_T/2\\pi m_ec^2)D=6\\sigma_T(1+\\sigma_w)\\gamma_w/m_ecL_w$. We now consider the term $rB$ that appears in equations (\\ref{dr1}), (\\ref{entropy}), and (\\ref{Lambda2}). Its form depends on whether the flow is in the ideal-MHD or the equipartition regime of the bubble interior. The ideal-MHD case corresponds to setting ${\\mathcal{E}}^\\prime = 0$ in equation (\\ref{EB}), which implies $rB = r{\\mathcal{E}}/\\beta\\propto 1/\\beta$ (by eq. [\\ref{Er}]). It is then straigthtforward to obtain from equations (\\ref{dr1}) and (\\ref{entropy}) (after also substituting $w=4p$ from eq. [\\ref{rel_eos}]) the following pair of coupled, first-order, ordinary differential equations for the variables $\\beta$ and $p$, which give the structure of the ideal-MHD sector of the PWB: \\begin{equation} {d\\beta\\over dr}= \\left[Gr^2{\\mathcal{E}}^2p^2+8\\gamma^2\\beta^2\\frac{p}{r}\\right ]\\left[4\\gamma^4\\beta(3\\beta^2-1)p -\\frac{3}{4\\pi}\\frac{{\\mathcal{E}}^2}{\\beta}\\right]^{-1}\\ , \\ \\ R_s\\le r \\le \\mathsf{min}\\{R_{\\rm eq}\\, ,\\, R_b\\}\\, ,\\label{beta_ideal} \\end{equation} \\begin{eqnarray}&& {dp\\over dr} = \\left [ 4 \\gamma^2\\beta^2 p - \\frac{{\\mathcal{E}}^2}{4\\pi \\gamma^2\\beta^2}\\right ] \\left[Gr^2{\\mathcal{E}}^2p^2 + 8\\gamma^2\\beta^2\\frac{p}{r} \\right ] \\left [ \\frac{3}{4\\pi}{\\mathcal{E}}^2 + 4\\gamma^4\\beta^2(1-3\\beta^2)p\\right ]^{-1} \\nonumber \\\\ && \\quad\\quad\\quad\\quad\\quad\\quad\\quad\\quad\\quad\\quad R_s\\le r \\le \\mathsf{min}\\{R_{\\rm eq}\\, ,\\, R_b\\}\\, , \\label{p_ideal} \\end{eqnarray} with ${\\mathcal{E}}$ given by equation (\\ref{Er}). For the nonideal (equipartition) regime, we combine equations (\\ref{delta}) and (\\ref{EB}) to obtain a quadratic equation for $rB$, whose relevant root is \\begin{equation}\\label{B_eq} rB_{\\rm eq} = \\frac{2\\beta}{1+\\beta^2}r{\\mathcal{E}} + \\frac{[8\\pi \\delta \\gamma^2 (1+\\beta^2)r^2 p-r^2{\\mathcal{E}}^2]^{1/2}}{\\gamma^2(1+\\beta^2)}\\ . \\end{equation} Using this relation as well as equations (\\ref{rel_eos}), (\\ref{Er}), and (\\ref{Lambda2}) in equations (\\ref{dr1}) and (\\ref{entropy}), it is once again possible to extract explicit differential equations for $\\beta$ and $p$: \\begin{eqnarray} && {d\\beta \\over dr}=\\left \\{ \\left [1+{\\delta (rB_{\\rm eq}-\\beta r{\\mathcal{E}})\\over \\gamma^2[(1+\\beta^2)rB_{\\rm eq}-2\\beta r{\\mathcal{E}}]}\\right ]G\\gamma^2\\beta (rB_{\\rm eq}-\\beta r{\\mathcal{E}})^2p^2 + \\nonumber \\right. \\\\ && \\left. +\\left [ 8\\beta +\\frac{2\\delta r {\\mathcal{E}}}{\\gamma^4[(1+\\beta^2)rB_{\\rm eq} -2\\beta r {\\mathcal{E}}]}\\right ]\\frac{p}{r}\\right \\} \\left \\{ 4 \\left [\\gamma^2 (3\\beta^2-1)-\\delta \\right ] p + \\right. \\nonumber \\\\ && \\left. +{(r^2{\\mathcal{E}}B_{\\rm eq}-8\\pi\\delta \\gamma^2\\beta r^2 p) \\left[4\\beta rB_{\\rm eq}-(1+3\\beta^2)r{\\mathcal{E}}\\right] \\over 2\\pi(1+\\beta^2)\\gamma^2[(1+\\beta^2)rB_{\\rm eq}-2\\beta r{\\mathcal{E}}]r^2} \\right \\}^{-1}\\, , \\nonumber \\\\ && \\quad\\quad\\quad\\quad\\quad\\quad\\quad\\quad\\quad\\quad \\mathsf{max}\\{R_s\\, ,\\, R_{\\rm eq}\\}\\le r \\le R_b\\, ,\\label{beta_eq} \\end{eqnarray} \\begin{eqnarray} &&{dp\\over dr} = \\left \\{1+\\frac{\\delta(rB_{\\rm eq}-\\beta r{\\mathcal{E}})}{\\gamma^2[(1+\\beta^2)rB_{\\rm eq}-2\\beta r{\\mathcal{E}}]}\\right \\}^{-1}\\times \\nonumber \\\\ && \\times \\left\\{\\left[{(\\beta r {\\mathcal{E}}-rB_{\\rm eq})(r^2{\\mathcal{E}}B_{\\rm eq}-8\\pi\\delta\\gamma^2\\beta r^2p)\\over 2\\pi(1+\\beta^2)\\gamma^2[(1+\\beta^2)rB_{\\rm eq}-2\\beta r {\\mathcal{E}}]r^2}-4\\gamma^2\\beta p\\right ] \\left({d\\beta \\over dr} \\right ) + {2\\delta(\\beta r{\\mathcal{E}}-rB_{\\rm eq})\\over \\gamma^2[(1+\\beta^2)rB_{\\rm eq}-2\\beta r{\\mathcal{E}}]} \\frac{p}{r} \\right \\}\\ , \\nonumber \\\\ && \\quad\\quad\\quad\\quad\\quad\\quad \\mathsf{max}\\{R_s\\, ,\\, R_{\\rm eq}\\}\\le r \\le R_b\\, ,\\label{p_eq} \\end{eqnarray} where the term $(d\\beta/dr)$ in equation (\\ref{p_eq}) is given by the expression (\\ref{beta_eq}). Equations (\\ref{beta_ideal})--(\\ref{p_eq}) are integrated over their respective validity domains subject to the boundary conditions \\begin{equation}\\label{BC} \\beta(R_s)=\\beta_{\\rm ps}\\, , \\quad\\quad p(R_s)=p_{\\rm ps}\\, , \\end{equation} where $\\beta_{\\rm ps}$ is given by equation (\\ref{beta_ps}) and the postshock pressure \\begin{equation}\\label{p_ps} p_{\\rm ps} = {\\sigma_w \\gamma_w^2\\beta_w^2n_w(R_s)m_ec^2 \\over 2\\delta_{\\rm ps} \\gamma_{\\rm ps}^2\\beta_{\\rm ps}^2} \\end{equation} (with $n_w(r)$ given by eq. [\\ref{Lsigma}]) is similarly obtained from the wind-shock jump conditions (see GK). The value of $R_s$, where the boundary conditions (\\ref{BC}) are imposed, is not known a priori and must be determined from an additional constraint. This can be provided by requiring global particle conservation: for a bubble considered at time $t$ after the supranova explosion, the total number of particles within the radius $R_b(t)$ [which consists of the unshocked wind at $rR_s(t)$] is equal to the total number of particles injected by the central neutron star over the time $t$. The pair injection rate at the source is given by \\begin{equation}\\label{Ndot} \\dot{N}= {L_w \\over (1+\\sigma_w)\\gamma_w m_e c^2}\\, , \\end{equation} and hence the total number of particles within $R_b$ at time $t$ is $N(t)=\\dot{N}t$. We approximate $t\\approx R_b/\\beta_bc$, which should be accurate to within a factor of order 1 (for example, $t=1.5R_b/\\beta_bc$ in the case of an adiabatic bubble, with the numerical coefficient decreasing in the presence of cooling; see \\cite{RC84}). The number of particles within the volume occupied by the unshocked wind is thus \\begin{equation}\\label{N_1} N(rv_b$ everywhere within this region and that the ejecta speed was $1$ in this case), the nonideal regime becomes progressively smaller with decreasing $\\sigma_w$. This is a direct consequence of equation (\\ref{delta_ps}), which indicates that $\\delta_{\\rm ps}$ scales approximately linearly with $\\sigma_w$: the lower the value of $\\delta_{\\rm ps}$, the longer it will take for $p_B/p$ in the postshock flow to rise above 1 (our condition for the termination of the ideal-MHD regime). The relative extent of the ideal-MHD region is larger for the fixed-$\\beta(R_b)$ solutions, so much so that the solutions of this type with the two lowest values of $\\sigma_w$ contain no nonideal-MHD zone. This can be understood from the systematically higher values of $\\beta(\\tilde{r})$ (see Figs. \\ref{fig2} and \\ref{fig3}) and correspondingly lower values of $\\tilde{B}(\\tilde{r}) \\propto 1/\\tilde{r}\\beta(\\tilde{r})$ and hence of $p_B(\\tilde{r})/p(\\tilde{r})$ in the fixed-$\\beta(R_b)$ solutions in comparison with their $N$-conserving counterparts. Figures \\ref{fig2} and \\ref{fig3} exhibit the $\\tilde{r}$ dependence of various quantities of interest in PWB solutions obtained by imposing the $N$-conservation and fixed-$\\beta(R_b)$ constraints, respectively, for the same 4 values of the wind magnetization parameter as in Figure \\ref{fig1} (spanning 3 orders of magnitude in $\\sigma_w$) and for 5 values of the cooling parameter (spanning 4 orders of magnitude in $a_1$). If $\\beta_{b,-1}$ and $\\Delta E_{53}$ are both set equal to 1, then the chosen values of $a_1$ ($=0.45,\\, 5.04,\\, 45.3,\\, 504$, and 4539) correspond to SMNS spin-down times $t_{\\rm sd}=100,\\, 30,\\, 10,\\, 3$ and $1\\ {\\rm yr}$, respectively. Bubbles with $a_1$ near the lower end of this range resemble adiabatic PWBs, whereas configurations with $a_1$ near the upper range are highly radiative. Radio pulsars have inferred surface magnetic fields in the range $\\sim 10^{12}-10^{13}\\ {\\rm G}$, so, by equation (\\ref{t_sd}), a variation of roughly two orders of magnitude in the value of $t_{\\rm sd}$ is naturally expected. The upper panels in Figures \\ref{fig2} and \\ref{fig3} show the behavior of $\\beta(\\tilde{r})$. The {\\it dashed}\\/ lines in these panels indicate the postshock ($\\beta_{\\rm ps}$) and outer-boundary ($\\beta_b$) speeds, whereas the {\\it dash-dotted}\\/ curves depict the purely adiabatic ($a_1=0$) solutions. The displayed results confirm that $\\beta(R_b)<\\beta_b$ in the $N$-conserving solutions. The greatest discrepancy between $\\beta(R_b)$ and the actual speed ($\\beta_b$) of the outer boundary occurs for $\\sigma_w=1$, in which case $\\beta(R_b)/\\beta_b$ decreases from $\\sim 0.13$ to $\\sim 0.0063$ as $a_1$ increases from 0.1 to $10^3$. We consider this to be a tolerable discrepancy, given (as we already noted in \\S \\ref{structure}) that the flow near $\\tilde{r}=1$ is highly subsonic. The $\\beta(\\tilde{r})$ curves further demonstrate that the values of $\\beta$ in the $N$-conserving solutions are lower than those in the corresponding $\\beta(R_b)=\\beta_b$ solutions also for all other values of $\\tilde{r}$ where the respective solutions overlap. The lower values of $\\beta$ lead (on account of the particle flux-conservation relation [\\ref{particle}]) to systematically higher values of $n$ in the $N$-conserving solutions (see the second row of panels in Figs. \\ref{fig2} and \\ref{fig3}). The radial profiles of $n$ are nearly flat for low values of the cooling parameter, but when radiative effects are important and contribute to the compression, $n(\\tilde{r})$ rises monotonically between $\\tilde{r}_s$ and $\\tilde{r}=1$ with a slope that is steeper (particularly in the nonideal-MHD regime) the larger the value of $a_1$. The third row of panels in Figures \\ref{fig2} and \\ref{fig3} displays $\\gamma_e(\\tilde{r})/\\gamma_w$. The {\\it dashed}\\/ line marks the postshock value of this quantity, and the {\\it dash-dotted}\\/ curves again represent the purely adiabatic case. The random Lorentz factor declines monotonically with $\\tilde{r}$ in the $\\sigma_w=1$ and 0.1 solutions, but at lower values of $\\sigma_w$ (for which the electromagnetic pressure contribution is small) and so long as cooling remains relatively unimportant ($a_1\\lesssim 10$), $\\gamma_e(\\tilde{r})$ initially increases behind the wind shock to make it possible to attain the necessary total pressure. The minimum value of $\\gamma_e(\\tilde{r})$ (reached at $\\tilde{r}=1$) decreases (as expected) with increasing $a_1$ and is lower for the $N$-conserving solutions than for the fixed-$\\beta(R_b)$ ones. This difference can be understood from the fact that the particle density near the outer boundary is lower in the latter case [corresponding to a higher value of $\\beta(\\tilde{r}=1)$], so $\\gamma_e$ must be larger to bring the pressure up to its requisite value. Of particular relevance to the evolution of GRB afterglows is the behavior of the thermal and electromagnetic pressures. The radial profiles of $p$ (normalized by $p_1$) are shown in the fourth row of panels in Figures \\ref{fig2} and \\ref{fig3}. It is seen that the behavior of $p(\\tilde{r})$ varies with the parameter choices and that the details depend on the nature of the constraint imposed on the solution. For $\\sigma_w=1$, the curves decline monotonically when cooling is relatively unimportant but increase monotonically at high values of $a_1$. For $\\sigma_w=0.1$, the curves decrease monotonically for all plotted values of the cooling parametr, whereas for $\\sigma_w=10^{-3}$ in the fixed-$\\beta(R_b)$ solution they increase monotonically for all exhibited values of $a_1$. In the remaining cases, the curves increase with $\\tilde{r}$ near the inner boundary of the bubble and decrease near its outer boundary. Since $p \\propto \\gamma_e n$, one can understand the shape of the curves by comparing them with the corresponding curves in the second and third rows of panels. One finds that, quite generally, the behavior of $p$ is dominated by that of $\\gamma_e$ for low values of $a_1$, but that the influence of the density variations becomes progressively more important as the cooling parameter increases. The effect of the electromagnetic fields can be inferred from the behavior of the variable $\\psi$ and of the ratio ${\\mathcal{E}}^\\prime/B^\\prime$, shown in the bottom two rows of panels in Figures \\ref{fig2} and \\ref{fig3}. Note that one can write $\\psi = \\delta + {\\mathcal{E}}^\\prime B^\\prime/4\\pi p$ (see eq. [\\ref{psi}]). Immediately behind the wind shock (in which, by construction, ${\\mathcal{E}}^\\prime=0$), $\\psi$ is equal to $\\delta_{\\rm ps}$ (given by eq. [\\ref{delta_ps}]). If $\\delta_{\\rm ps}\\ge 1$ then the nonideal-MHD regime starts right there: beyond that point, $\\delta$ remains fixed at its postshock value but $\\psi$ and ${\\mathcal{E}}^\\prime/B^\\prime$ increase with $\\tilde{r}$. If $\\delta_{\\rm ps}< 1$, then the evolution proceeds under ideal-MHD conditions (${\\mathcal{E}}^\\prime=0$), with $\\psi$ continuing to coincide with $\\delta$ (which in this regime equals $p_B/p$) up to the point where it reaches 1. This point marks the end of the ideal-MHD regime and corresponds to $\\tilde{r}_{\\rm eq}$. Beyond $\\tilde{r}_{\\rm eq}$, $\\psi$ continues to increase monotonically with $\\tilde{r}$ (and now so does also ${\\mathcal{E}}^\\prime/B^\\prime$), but $\\delta$ remains fixed at 1 (see eq. [\\ref{delta}]). The figures show that, as expected, the electromagnetic contribution to the total pressure becomes progressively larger as $\\sigma_w$ increases, and they further demonstrate that ${\\mathcal{E}}^\\prime/B^\\prime$ exhibits a similar trend. It is also seen that cooling enhances the relative importance of the electromagnetic fields, which can be understood from the fact that it reduces the magnitude of the thermal pressure component. The overall behavior of $\\psi$ and ${\\mathcal{E}}^\\prime/B^\\prime$ does not appear to depend strongly on the choice of constraint under which the solution is obtained, although the corresponding curves in the two figures differ in their details (see also the related discussion in connection with Fig. \\ref{fig1} above). As we show in \\S \\ref{parameters}, the variable $\\psi$ plays a key role in the modeling of relativistic shocks that propagate inside a PWB. ", + "conclusions": "\\label{conclusions} We propose to identify the environment into which afterglow-emitting shocks in at least some GRB sources propagate with pulsar-wind bubbles. Our results can be summarized as follows: \\begin{itemize} \\item PWBs provide a natural resolution of the apparent difficulty of accounting for the high electron and magnetic energy fractions ($\\epsilon_e$ and $\\epsilon_B$, respectively) inferred in a number of afterglow sources. This is because pulsar winds are expected to have a significant $e^\\pm$ component and to be highly magnetized. If high values of $\\epsilon_e$ in fact prove to occur commonly in afterglow sources, then this would strengthen the case for a simple, ``universal'' explanation of this type. \\item An association of PWBs with GRBs is expected under several GRB formation scenarios, including the collapse of a massive star. In light of suggestive evidence that many of the afterglows observed to date may have a massive stellar progenitor, we have concentrated on this case. In particular, we considered the supranova scenario of VS, in which an intense pulsar-type wind from the GRB progenitor is a key ingredient of the hypothesized evolution. In this picture, the ejection of a highly energetic, ultrarelativistic pulsar wind is predicted to follow the supernova explosion and to last anywhere from several months to several years until the central object collapses to form a black hole, thereby triggering the burst. Recent detections of X-ray features in several GRB sources have been interpreted as providing strong support for this scenario. \\item To assess the implications of a PWB environment to afterglow sources in the context of the supranova scenario, we have constructed a simple, steady-state model of the early-time structure of a plerionic supernova remnant. We have been guided by Atoyan's (1999) spectral modeling of the Crab, which yielded a lower initial wind Lorentz factor and a higher initial pulsar rotation rate than in previous estimates, and by other recent results on the Crab and Vela synchrotron nebulae, from which we inferred a plausible range of the wind magnetization parameter $\\sigma_w$ ($\\sim 10^{-3}-1$). In contradistinction to previous models of the structure of plerionic SNRs, we have replaced the assumption that ideal MHD applies throughout the PWB with the postulate that the electromagnetic-to-thermal pressure ratio in the bubble remains constant after it increases to $\\sim 1$. We have also explicitly incorporated synchrotron-radiation cooling. Although our solutions do not provide an exact representation of radiative (and thus intrinsically time-dependent) PWBs, we have verified that they generally do not depend on the detailed approximations that are adopted and are essentially characterized by $\\sigma_w$ and by a second parameter that measures the relative importance of radiative cooling within the bubble. It would be of interest to further develop this model and to investigate the possibility that it can be applied both to young radio pulsars and to GRB progenitors as members of the same general class of rapidly rotating and strongly magnetized neutron stars. \\item In view of the ``hot'' (relativistic) equation of state and high magnetization of the shocked wind, the effective hydrogen number density that determines the properties of a relativistic afterglow-emitting shock is given by $n_{\\rm H, equiv}= [4p+(B^\\prime+{\\mathcal{E}}^\\prime)^2/4\\pi]/ m_p c^2$, where $B^\\prime$ and ${\\mathcal{E}}^\\prime$ are, respectively, the comoving magnetic and electric fields and $p$ is the particle pressure. For plausible values of the cooling parameter (and independent of the value of $\\sigma_w$), the derived values of $n_{\\rm H, equiv}$ span the density range inferred from spectral modeling of GRB afterglows. An interesting feature of the solutions is the predicted radial variation of $n_{\\rm H, equiv}$ within the bubble, which can mimic either a uniform-ISM or a stellar-wind environment, but which in general exhibits a more diverse behavior. Among other things, this model makes it possible to understand how a GRB with a massive progenitor can produce an afterglow that shows no evidence of a stellar-wind or a high-density environment. \\item We have examined the dependence of the characteristic synchrotron spectral quantities in an afterglow-emitting shock that propagates inside a PWB on the bubble parameters and related them to the standard expressions derived under the assumption of a uniform-ISM environment. We found that, under typical circumstances, the standard expressions remain roughly applicable if one substitutes for $\\epsilon_e$, $\\epsilon_B$, and $n_{\\rm H}$ their ``equivalent'' PWB expressions. We noted, however, that the parameter scaling laws would change in strongly radiative bubbles: these differences might be detectable in objects with high inferred ambient densities. \\item Finally, we considered the possible observational manifestations of the dense supranova shell that surrounds the PWB in this picture. In particular, we discussed how the X-ray emission features detected in objects like GRB 991216 may be interpreted in the context of a supranova-generated PWB. We concluded that both the X-ray features and the afterglow emission could be explained by this model if the PWB were elongated, and we argued that such a shape might be brought about by anisotropic mass outflows from the GRB progenitor star. \\end{itemize}" + }, + "0112/astro-ph0112197_arXiv.txt": { + "abstract": "We present dynamically consistent solutions for hot accretion onto unmagnetized, rotating white dwarfs (WDs) in five quiescent dwarf novae. The measured WD rotation rates (and other system parameters) in RX And, SS Cyg, U Gem, VW Hyi and WZ Sge imply spindown of the WD by an extended hot flow emitting most of its X-rays in the vicinity of the stellar surface. In general, energy advection is absent and the flow is stable to convection and hydrodynamical outflows. In rapidly rotating systems, the X-ray luminosity provides only an upper limit on the quiescent accretion rate because of substantial stellar spindown luminosity. We suggest that the presence of hot flows in quiescent dwarf novae may limit the long-term WD rotation rates to significantly sub-Keplerian values. ", + "introduction": "Dwarf novae (DN) are accreting binary star systems with a white dwarf (WD) primary fed by a main-sequence donor via Roche-lobe overflow. They represent a subclass of cataclysmic variables (CVs) and experience semi-regular, luminous outbursts, during which accretion onto the WD proceeds at a high rate. Most of the time, however, DN are in quiescence --- a phase during which the accretion rate onto the WD (and the system luminosity) is considerably reduced (see Warner 1995 for a review). According to the disk instability model (DIM), such a behavior arises because most of the mass transferred by the companion star builds up in an unsteady disk during quiescence (Cannizzo 1993; Lasota 2001). Quiescent DN are hard X-ray sources with typical luminosities of $\\sim 10^{30}-10^{32}$~erg~s$^{-1}$ (see, e.g., C\\'ordova \\& Mason 1983; Patterson \\& Raymond 1985). Spectral fits to this X-ray emission suggest a Bremsstrahlung origin from gas with temperatures $\\sim 2-20$~keV (Patterson \\& Raymond 1985; Eracleous et al. 1991; Belloni et al. 1991; Yoshida et al. 1992; Mukai \\& Shiokawa 1993). This X-ray emission is commonly attributed to the boundary layer (BL) at the interface between the WD and the thin accretion disk around it. At low accretion rates ($\\lsim 10^{16}$~g~s$^{-1}$, typical of quiescent DN), the gas in the BL is hot and optically thin, hence it is a substantial source of hard X-ray emission (Pringle \\& Savonije 1979; Tylenda 1981; Patterson \\& Raymond 1985). Detailed calculations by Narayan \\& Popham (1993) show that the optically-thin BLs of disk-accreting WDs can also be radially extended (of order a WD radius) and that energy advection is an important element of their internal structure. On the other hand, Meyer \\& Meyer-Hofmeister (1994) presented two arguments in favor of a more extended, inner hot flow structure in quiescent DN. They noted that: (i) the standard DIM predicts accretion rates at the WD surface which are smaller than those inferred from quiescent X-ray luminosities by typically more than one order of magnitude and (ii) the presence of an extended, low-density hot flow in the WD vicinity may explain the observed delay (of $\\sim 0.5-1$ day) in the rise to outburst of the EUV light relative to the optical light (e.g., Mauche, Mattei \\& Bateson 2001, but see Smak 1998). Our current understanding of the structure and properties of quiescent disks is rather limited (Menou 2002), so studying the possible presence of an extended, hot flow in quiescent DN is certainly worthwhile. Recently, Medvedev \\& Narayan (2001a) discovered solutions for hot accretion onto unmagnetized, rotating compact stars. They found that at large stellar spin rates, dissipation in the hot flow is dominated by stellar braking in a ``hot settling flow'' (HSF) configuration, whereas in the opposite limit, the flow reduces to a conventional advection-dominated accretion flow (ADAF; Narayan \\& Yi 1994; 1995). Both flows exist at relatively low accretion rates ($\\la10^{-2}$ of Eddington). Medvedev \\& Narayan (2001b) have also addressed the thermal stability of their (cooling-dominated) settling solution, and found that it is most likely stable, once { turbulent} thermal conduction is accounted for. The work of Medvedev \\& Narayan was largely focused on accretion onto neutron stars, but the HSF solutions are also valid for WD accretion. This is important because, in general, much more is known about WDs in CVs than about neutron stars in close binary systems. In particular, the rotation rates of several WDs in DN have been measured during the last few years thanks to {\\it HST} spectroscopy (see Sion 1999 for a review). In this {\\it Letter}, we apply the model of Medvedev \\& Narayan to WDs in quiescent DN by constructing numerical solutions for hot accretion in five systems with relatively well known system parameters. ", + "conclusions": "We constructed numerical models for hot accretion onto unmagnetized, rotating WDs in five quiescent DN. Our solutions are a significant improvement over previous work on hot accretion in this context (Katz 1977; Kylafis \\& Lamb 1982; Mahasena \\& Osaki 1999; Menou 2000) in that both the WD rotation and the viscous nature of the flow are accounted for in the present case. This makes these solutions plausible modes of accretion in quiescent DN. The outer boundary conditions chosen are somewhat artificial and constitute a significant shortcoming of the present study. It is unclear exactly how the transition from an inner hot flow to an outer thin disk (known to be present) would occur. Strong boundary effects may be expected for a hot flow with a limited radial extent. Nonetheless, our work has the virtue of isolating the main properties of the hot flow, independently of the disk. It will be important in the future to study the structure of a hot flow with a smaller radial extent and more realistic boundary conditions. For the non-relativistic gas considered here and non-dominant magnetic fields, we expect rather large values of $\\gamma$ and therefore flows which are not subject to outflows (Blandford \\& Begelman 1999) or convection (Narayan et al. 2000; Quataert \\& Gruzinov 2000). Menou (2000) proposed that energy advection in a hot flow could be important in powering a dominant EUV emission component that could explain the strong He~II emission lines observed in many quiescent DN. Since the hot flow solutions presented here lack energy advection (for $\\gamma$'s near 5/3), this possibility seems unlikely. While a star is spun up by accretion via a disk boundary layer (except when rotating near breakup; Popham \\& Narayan 1991; Paczynski 1991), it is spun down when accreting via a hot flow like an ADAF, a HSF or the hot flow solutions presented here (except at very low spin rates; Medvedev \\& Narayan 2001a). This property has potentially important consequences for the spin history of WDs in DN. Sion (1999) emphasized that in the standard picture WDs should rotate near break up in DN, but observations suggest that rapidly spinning WDs are rare. If a WD is spun down via hot accretion during quiescence, an equilibrium at substantially sub-Keplerian rotation rates may be expected from the balance between spin-up in outburst and spin-down in quiescence. This possibility requires further study with more appropriate boundary conditions. We note that other mechanisms could contribute to spinning down single WDs (Spruit 1998) and WDs in DN (and thus get rid of the angular momentum acquired during accretion), including angular momentum losses induced by coupling to an expanding envelope during nova explosions (Livio \\& Pringle 1998; Sion et al. 2001). X-ray eclipses in several quiescent DN indicate that the emission originates from the vicinity of the WD (Wood et al. 1995; Mukai et al. 1997; Pratt et al. 1999; Ramsay et al. 2001). Without a detailed comparison, it is unclear whether the hot flow solutions presented here satisfy the existing observational constraints or not. Such a comparison should take into account the spectral coverage of the various X-ray instruments used and it probably requires a reliable boundary layer model taking into account the role of heat conduction in potentially modifying its emission properties (which is beyond the scope of the present paper). It is certainly encouraging that X-ray emission is highly concentrated near the stellar surface in our models, but it is presently unclear whether these models will be able to account for the available X-ray eclipse data or not. We note, for instance, that the X-ray eclipse of OY Car observed by Ramsay et al. (2001) suggest an X-ray flux approaching zero at eclipse minimum and a small vertical extent for the X-ray emitting region. These properties may not be easily reconciled with the hot flow models presented here. In the future, it will also be possible to probe the structure of hot flows in (even non-eclipsing) quiescent DN with detailed X-ray spectroscopic diagnostics, as illustrated in Menou, Perna \\& Raymond (2001; see also Narayan \\& Raymond 1999)." + }, + "0112/astro-ph0112368_arXiv.txt": { + "abstract": "HI data cubes are sources of unique information on interstellar turbulence. Doppler shifts due to supersonic motions contain information on turbulent velocity field which is otherwise difficult to obtain. However, the problem of separation of velocity and density fluctuations within HI data cubes is far from being trivial. Analytical description of the emissivity statistics of channel maps (velocity slices) in Lazarian \\& Pogosyan (2000) showed that the relative contribution of the density and velocity fluctuations depends on the thickness of the velocity slice. In particular, power-law assymptotics of the emissivity fluctuations change when the dispersion of the velocity at the scale under study becomes of the order of the velocity slice thickness (integrated width of the channel map). These results are the foundations of the Velocity-Channel Analysis (VCA) technique which allows to determine velocity and density statistics using 21-cm data cubes. The VCA has been successfully tested using data cubes obtained via compressible magnetohydrodynamic simulations and applied to Galactic and Magellanic Clouds data. As a tool it has become much more sophisticated recently when effects of absorption were accounted for. The systematic studies of vast 21-cm data sets to correlate the variations in the turbulence statistics with the astrophysical activity is on the agenda. This should allow to determine the interstellar energy injection mechanisms. Going beyond the VCA, we discuss other tools, namely, genus and anisotropy analysis. The first characterises the topology of HI, while the second provides magnetic field directions. We show a few applications of these new tools to HI data and MHD simulations. ", + "introduction": "Atomic hydrogen is an important component of the interstellar media of spiral galaxies (McKee \\& Ostriker 1977) and much efforts have been devoted to its studies (see this volume). From the point of view of fluid mechanics, HI, as well as other components of interstellar medium, is characterised by huge Reynolds numbers, $Re$, which is the ratio of the eddy turnover time of a parcel of gas to the time required for viscous forces to slow it appreciably. For $Re\\gg 100$ we expect gas to be turbulent and this is exactly what we observe in HI (for HI $Re\\sim 10^8$). Statistical description is a nearly indispensable strategy when dealing with turbulence. The big advantage of statistical techniques is that they extract underlying regularities of the flow and reject incidental details. Attempts to study interstellar turbulence with statistical tools date as far back as the 1950s (see Horner 1951, Kampe de Feriet 1955, Munch 1958, Wilson et al. 1959) and various directions of research achieved various degree of success (see reviews by Kaplan \\& Pickelner 1970, Dickman 1985, Lazarian 1992, Armstrong, Rickett \\& Spangler 1995). Studies of turbulence statistics of ionized media were successful (see Spangler \\& Gwinn 1990) and provided the information of the statistics of plasma density at scales $10^{8}$-$10^{15}$~cm. This research profited a lot from clear understanding of processes of scintillations and scattering achieved by theorists (see Narayan \\& Goodman 1989). At the same time the intrinsic limitations of the scincillations technique are due to the limited number of sampling directions, relevance only to ionized gas at extremely small scales, and impossibility of getting velocity (the most important!) statistics directly. ``Seeing through dust'' at 21 cm potentially provides an enormous leap for turbulence studies. First of all, statistics of HI line carries the velocity information. Secondly, the statistical samples are extremely rich and not limited to discrete directions. Thirdly, HI emission allows to study turbulence at large scales which have direct relation to the scales of star formation and energy injection. Deficiencies in the theoretical description has been, to our mind, the major impediment to studies of turbulence using 21-cm data cubes. In particular, Crovisier \\& Dickey (1983) Green (1993) measured the spectrum of 21-cm intensity fluctuations, but it was unclear what those spectra mean. Sally Oey (this volume) identifies several known problems of studying HI turbulence. They are ({\\it a}) integration of fluctuations along lines of sight, ({\\it b}) contribution of velocity and density to the emissivity statistics and ({\\it c}) effects of absorption. The problems related to ({\\it a}) were dealt in Lazarian (1995), while the issues ({\\it b}) and ({\\it c}) happen to be more formidable and have been dealt with only recently (Lazarian \\& Pogosyan 2000, 2002). Earlier reviews dealing with HI studies include Lazarian (1999a,b). In this volume one can read more about turbulence in Cho, Lazarian \\& Yan and Vazquez-Semadeni papers. ", + "conclusions": "For the first time ever we have an adequate theoretical description of spectral line statistics in the presence of both velocity and density fluctuations. This must be exploited to get adequate description of interstellar turbulence. HI data cubes present a great opportunity for such a study. To get a better understading of interstellar dynamics we propose using complementary tools like {\\it genus} and {\\it anisotropy analysis} and correlating the statistics of HI with that of molecular species and ions. {\\bf Acknowledgement}. AL acknowledges the support by the grant NSF AST-0125544." + }, + "0112/astro-ph0112532_arXiv.txt": { + "abstract": "{ We observed an anomalously outbursting state of SU~UMa which occurred in 1992. Time-resolved photometry revealed the presence of signals with a period of 0.0832$\\pm$0.0019~d, which is 3.6$\\sigma$ longer than the orbital period (0.07635~d) of this system. We attributed this signal to superhumps, based on its deviation from the orbital period and its characteristic profile. During this anomalous state of SU~UMa, normal outbursts were almost suppressed, in spite of relatively regular occurrences of superoutbursts. We consider that an ensuing tidally unstable state following the preceding superoutburst can be a viable mechanism to effectively suppress normal outbursts, resulting in an anomalously outbursting state. ", + "introduction": "SU~UMa is the prototype of SU~UMa-type dwarf novae (for a recent review of SU~UMa-type stars and their observational properties, see \\citealt{war95suuma}). Although SU~UMa shows typical superoutbursts and associated superhumps as in other SU~UMa-type dwarf novae \\citep{uda90suuma}, the star is also known to sometimes show anomalous states lacking superoutbursts, or even normal outbursts (cf. \\citealt{ros00suuma}). According to \\citet{ros00suuma}, the period of 1980--1983 was the most remarkable, when SU~UMa almost completely stopped outbursting. Several instances have been known, that SU~UMa showed a lesser degree of anomalously outbursting state. February, 1992 was another such period (cf. \\citealt{ros00suuma}), when SU~UMa ceased to show normal outbursts. During that period, the mean brightness of SU~UMa was observed brighter than the averaged quiescent level. ", + "conclusions": "The period analysis strongly supports that presence of the period of 0.0832$\\pm$0.0019~d. The error of the period was estimated using the application of Lafler-Kinman class of methods by \\citet{fer89error}. Although the error of period estimation is rather large due to the limited length of a single-night baseline, the period is 3.6$\\sigma$ longer than the orbital period, and is 2.1$\\sigma$ longer than the superhump period by \\citet{uda90suuma}. Since superhump periods can vary to a considerable extent (e.g. \\citealt{kat01aqeri}), the present periodicity, which is significantly longer than the orbital period, is more regarded as a variety of superhumps, rather than the one reflecting the orbital period. \\begin{figure} \\includegraphics[angle=0,height=6cm]{H3348F4.eps} \\caption{Phase-averaged light curve (1992 February 17).} \\label{fig:phase} \\end{figure} The hump profile shown in Fig. \\ref{fig:phase}. A rapid rise and slower fade, and the presence of a secondary hump at phase 0.6 (cf. \\citet{uda90suuma} for the discussion of secondary superhumps during a superoutburst), are also very characteristic of superhumps. Such appearance of the superhump signal may be related to the anomalous state of SU~UMa. \\citet{ros00suuma} argued that the complete cessation of outburst in the period of 1980--1983 was probably caused by a strong variation of mass-transfer rates ($\\dot{M}$). This interpretation, however, does not seem to adequately explain the 1992 anomalous state, since the interval of superoutbursts (labeled S020 on JD 2448784 and S021 on JD 2449047 in \\citet{ros00suuma}) just before and after this anomalous state was 263 d, which is only slightly longer than the typical value of this object (see Fig. 5 of \\citet{ros00suuma}). This interval is also a very typical value as seen in other SU~UMa-type dwarf novae \\citep{nog97sxlmi}. Since the length of a supercycle is primarily governed by the transferred angular momentum from the secondary star \\citep{osa89suuma}, the interval of successive superoutbursts is roughly inversely proportional to $\\dot{M}$ \\citep{ich94cycle}. This suggests that $\\dot{M}$ was relatively normal during this anomalous state, and that normal outbursts may have been somehow suppressed even under the condition of a usual $\\dot{M}$. There are some known SU~UMa-type dwarf novae which show a permanently, or temporarily, reduced frequency of normal outbursts, in contrast to their high frequency of superoutbursts (V503~Cyg: \\citealt{har95v503cyg}, Ishioka et al. in preparation; V1113~Cyg: \\citealt{kat01v1113cyg}). The presence of active and inactive phases in V1113~Cyg, in terms of the frequency of normal outbursts, while maintaining the supercycle length, seems to require an unknown mechanism which effectively suppresses normal outbursts \\citep{kat01v1113cyg}. The present discovery of a signal, which can be attributed to superhumps, during a similarly anomalously outbursting state of SU~UMa provides an additional clue to this phenomenon. SU~UMa during anomalous outbursting states may more or less resemble permanent superhumpers (cf. \\citealt{osa96review}), which do not show strong dwarf nova-type outbursts, but show superhumps. The accretion disk in permanent superhumpers is believed to be hot enough to suppress usual thermal instability, which is responsible for normal outbursts, but is tidally unstable, giving rise to permanent superhumps \\citep{osa96review}. Altough it is not yet clear whether such a condition is met during this anomalously outbursting state of SU~UMa, or whether a similar explanation can be applicable to SU~UMa-type dwarf novae with strongly variable activities, the possibility of an ensuing tidally unstable state following the preceding superoutburst can be a viable mechanism to effectively suppress normal outbursts. Since such anomalous states are known to be rather infrequent, intensive target-of-opportunity observations are strongly encouraged when future occurrence of such a state is recognized. \\vskip 3mm This work is partly supported by a grant-in aid (13640239) from the Japanese Ministry of Education, Culture, Sports, Science and Technology." + }, + "0112/astro-ph0112062_arXiv.txt": { + "abstract": "I discuss recent seminal work on the LARCS dataset: a panoramic study of rich clusters of galaxies at $z\\sim0.12$. The importance of observing beyond the cluster core is illustrated by exploiting these data to examine colour gradients across the clusters. ", + "introduction": "The Las Campanas / AAT Rich Cluster Survey (LARCS; O'Hely et al. 1998; Pimbblet et al. 2001a; 2001b; Pimbblet 2001) is a long-term project to study a statistically-reliable sample of 21 of the most luminous X-ray clusters at intermediate redshifts ($z=0.07$--0.16) in the southern hemisphere. The photometric imaging of these clusters comprises homogeneous, two degree wide $B$ and $R$ band observations taken at Las Campanas Observatory. These data permit tracing of photometric variations in cluster members out to large radii (typically $\\sim12$ Mpc at $z\\sim0.12$). Galaxies selected from these data are being observed in an on-going spectroscopic follow-up (Pimbblet 2001) with the 2dF spectrograph on the Anglo-Australian Telescope. ", + "conclusions": "To conclude, wide-field observations are paramount to the understanding of cluster formation and evolution and to trace radial variations in the cluster population (i.e. the CMR is readily observed to beyond 8 Mpc in many LARCS clusters; Pimbblet 2001). Although some recent work on a single high redshift cluster has already been undertaken by Kodama et al. (2001), much further work (photometric and spectroscopic) is urgently required to comparatively examine radial trends within a larger, well-defined sample of higher redshift clusters. Such observations would allow the testing of the prediction that these trends will be more strongly pronounced at higher redshifts (Pimbblet et al. 2001b)." + }, + "0112/astro-ph0112548_arXiv.txt": { + "abstract": "We compute the profiles of resonance doublet lines ($S_{1/2}-P_{1/2,3/2}$) formed in bipolar winds with velocity greater than the doublet separation in symbiotic stars. Particular attention has been paid on the doublet line ratio, where an essential role is played by the conversion of the short wavelength component arising from the $S_{1/2}-P_{3/2}$ transition into the long wavelength component for the transition $S_{1/2}-P_{1/2}$. We adopted a Monte Carlo technique and the Sobolev approximation. Our bipolar winds take the form of a cone and are characterized by the terminal wind velocity, the mass loss rate and the opening angle of the cone. When an observer is in the polar direction and the Sobolev optical depth $\\tau_{Sob}\\simeq 1$, we mainly obtain profiles with inverted flux line ratios, where the short wavelength component is weaker than the long wavelength component. When an observer is in the equatorial direction, we find that the profiles are characterized by two broad components, where the long wavelength component is the broader and stronger of the two. We conclude that the profiles obtained in our model provide a qualitative understanding of broad profiles and inverted intensity ratios of the doublets in symbiotic stars. ", + "introduction": "Symbiotic stars are generally known to be interacting binaries consisting of a red giant (or a Mira-type variable) and a hot companion that is usually a white dwarf (e.g. Kenyon 1986). Most red giant components suffer a heavy mass loss in the form of a slow stellar wind with a typical terminal speed 10 $\\sim20 {\\rm\\ km\\ s^{-1}}$ that is comparable to the escape velocity of a giant (e.g. Schmid 1996). In contrast, many white dwarf systems including planetary nebulae are known to possess fast outflows with terminal speed $\\ge 1000\\rm\\ km\\ s^{-1}$. Similar fast winds are also known in some symbiotic stars including AG~Peg (Vogel \\& Nussbaumer 1994). The orbital elements of symbiotic stars are not well-constrained, but light curves often show that they possess a long period of several hundred days (e.g. Iben \\& Tutukov 1996). M\\\"urset \\& Schmid (1999) presented the relation between the spectral types of the cool giants and the orbital periods to reveal that almost all symbiotic stars are well-detached binary systems. Therefore the most important binary activity may be found from the interaction between the two different kind of winds. These two winds may collide forming a shocked region, which can be identified with X-ray observations and is consistent with the fact that several symbiotic stars are known to be X-ray sources (Girard \\& Wilson 1987, M\\\"urset, Wolff \\& Jordan 1997, Ezuka, Ishida \\& Makino 1998) In many symbiotic stars the Raman-scattered O~{\\tiny{VI}} $\\lambda \\lambda6827, 7088$ features exhibit multiple peak profiles and strong polarization accompanied by the polarization flip (Schmid 1989, Schmid \\& Schild 1994, Harries \\& Howarth 1996). Lee \\& Park (1999) proposed that these features can be explained by assuming that there is an accretion disk around the white dwarf formed through capture of the slow wind from the giant (e.g. Mastrodemos \\& Morris 1998). According to the theoretical modelling by Paczynski \\& Zytkow (1978), periodic hydrogen shell flashes may occur in a white dwarf with the accretion rate $10^{-11} - 10^{-7}\\rm\\ M_{\\odot}\\ yr^{-1}$, where each outburst may last for decades. With this eruption the radiative pressure will drive a stellar wind around a white dwarf, which may take a bipolar form subject to the circumstellar matter distribution \\cite{sok2}. This interpretation is interesting because most symbiotic stars with known nebular morphologies are bipolar \\cite{cor}. Currently it is very controversial whether bipolar planetary nebulae possess central binary systems (e.g. Soker 1998). It is, therefore, very interesting and important to investigate the line profiles that indicate the fast outflowing motion, from which we may find the physical properties associated with the bipolarity of the wind. Since the resonance doublets arise from the common electronic transitions $S_{1/2}-P_{1/2,3/2}$ and have the separation ranging from 500${\\rm\\ km\\ s^{-1}}$ (for C~{\\tiny{IV}}) to 1,650 ${\\rm\\ km\\ s^{-1}} $ (for O~{\\tiny{VI}}), these lines can be an excellent tool to investigate the outflowing hot wind around the white dwarf component. When the bulk velocity changes significantly in a region that is much smaller than the scale height of the physical quantities such as density, the radiative transfer can be described by the Sobolev approximation. In this case, the optical depth for a line photon is inversely proportional to the velocity gradient in the direction of the photon propagation (e.g. Sobolev 1947; Rybicki \\& Hummer 1978). In the case of resonance doublet lines, line photons arising from the $S_{1/2}-P_{3/2}$ transition will be resonantly scattered with $S_{1/2}-P_{1/2}$ transition by receding ions with the speed of the doublet separation, if the outflow is an accelerating wind with a speed larger than the doublet separation. We denote this type of scattering by `double scattering.' This double scattering converts line photons associated with the $S_{1/2}-P_{3/2}$ transition into line photons associated with the $S_{1/2}-P_{1/2}$ transition. This will change the intrinsic doublet line ratio, which is 2~:~1 in the optically thin limit and 1~:~1 in the optically thick limit. Olson (1982) investigated this problem and applied to stellar winds around O and B stars. His main concern was limited to the investigations of various P~Cygni profiles in spherical winds. \\textit{IUE} observations indicate that many objects including symbiotic stars and planetary nebulae show various doublet line ratios between 2~:~1 and 1~:~1 (Feibelman 1983, Schmid et al. 1999). The microphysical processes that lead to these various line ratios may be associated with the existence of dust and/or collisional de-excitation (Ahn \\& Lee 2002, in preparation). Michalitsianos et al. (1988) showed that some symbiotic stars including RX Pup and R Aqr show anomalous line ratios, in which the short wavelength component of the C~{\\tiny{IV}} doublet was observed to be weaker than the long wavelength component. Vogel \\& Nussbaumer (1994) showed that the symbiotic nova AG Peg exhibits broad He~{\\tiny{II}} emission lines that are formed in a fast wind. In this system, the resonance doublets N~{\\tiny{V}} and C~{\\tiny{IV}} exhibit inverted line ratios where the short wavelength component is weaker. The authors proposed that the line formation is strongly affected by the double scatterings, whereby significant fraction of photons are converted. Furthermore, the doublets did not display P~Cygni profiles. These facts imply that the outflowing motion is confined to specific directions (plausibly in the polar directions), which exclude the observer's line of sight. There have been many theoretical investigations on the P~Cygni profiles shown in the resonance lines of metal elements by adopting the Sobolev approximation. However, relatively little attention has been paid on the profiles observed outside the wind flowing direction. In this paper, we perform Sobolev Monte Carlo computations to obtain the profiles of resonance doublet lines formed in bipolar winds that may be present in symbiotic stars. In Section 2, we briefly describe the Sobolev theory and the basic atomic physics concerning the resonant doublet lines. We also present the kinematic stellar wind model adopted in this work and the Monte Carlo procedure. Our results are presented in Section 3. Finally, we summarise and discuss our results and observational implications in Section~4. ", + "conclusions": "In this paper we investigate the doublet resonance line formation in bipolar conic flows in symbiotic stars by adopting the Sobolev Monte Carlo method. We also pay special attention to the doublet line ratios. When an observer is in the polar direction, there appear two peaks at the line centres coinciding with those of the resonance doublets. This is because we chose the mass loss rate $\\dot M$ so that the Sobolev optical depth $\\tau_{Sob} \\simeq 1$ for most of the velocity space, and therefore the line flux of the short component is suppressed by a factor $e^{-\\tau_{Sob}} \\simeq 0.3$. The red part of the long wavelength peak does not suffer any suppression. This fact implies that the detailed peak profiles differ in a systematic way. The suppression pattern also repeats for continuum parts, where the weakest occurs blueward of the short wavelength peak. Michalitsianos et al. (1988) investigated the profiles of C~{\\tiny{IV}} $\\lambda \\lambda 1548, 1551$ in the symbiotic stars RX~Pup and R~Aqr, from which they found the weaker short wavelength component than the long wavelength component. They particularly noted that the profiles are characterized by significant suppression in the blue part of the short wavelength peak. The peak flux ratio is most sensitively dependent on the Sobolev optical depth, particularly when $\\tau_{Sob}\\simeq 1$ and is insensitive on the covering factor of the wind. In RX~Pup and R~Aqr, there may exist a scattering component outflowing with speed less than the doublet separation. In this case, only single scattering is operational, which can cause the deviations of the doublet line ratio from the optically thin limit. However, when double scattering is possible, much more various doublet line ratios can be produced. It is uncertain that there exists an outflow moving faster than the doublet separation of 500${\\rm\\ km\\ s^{-1}}$ in RX~Pup and R~Aqr. In order to find the concrete evidence of double scattering, we need to obtain spectra with sufficient quality to discern the continuum level which are contributed by the redistributed photons. When we observe in the equatorial direction, we find that the profiles are characterized by two broad components, where the long wavelength component is the broader and stronger of the two. This qualitative feature does not change even if we vary either the mass loss rate or the Sobolev optical depth $\\tau_{Sob}$ as long as $\\tau_{Sob}$ exceeds unity. However, when the observer is in the polar direction, various P~Cygni type profiles are obtained dependent on the mass loss rate. We believe that the bipolarity of the stellar wind is mainly responsible for these spectroscopic characteristics. It is quite interesting that the symbiotic nova AG Peg shows the profiles of N~{\\tiny{V}} and C~{\\tiny{IV}} in the hot wind phases similar to those obtained in this work. We strongly believe that the hot wind in AG Peg takes the bipolar form with the speed exceeding $10^3{\\rm\\ km\\ s^{-1}}$. Kenny, Taylor \\& Seaquist (1991) used the Very Large Array (VLA) in order to investigate the morphological structure of AG Peg. They found a bipolar outer nebula extending $\\sim 40{\\rm\\ ''}$. They also found a bipolar enhancement in the inner nebula having a subarcsecond extent, which is consistent with the existence of a bipolar wind. However, it should also be noted that Nussbaumer, Schmutz \\& Vogel (1995) reported a P~Cygni type absorption in N~{\\tiny{V}} in AG~Peg using the {\\it Hubble Space Telescope}. Most symbiotic stars are not optically resolved and therefore their morphology is very uncertain. In planetary nebulae, it is still highly controversial whether the binarity of the central star system can be linked to the bipolarity of the nebula (e.g. Morris 1987, Livio \\& Soker 1988, Soker 1998). A significant input should be provided from the morphological study of symbiotic stars. In this respect, it is proposed that the resonance doublets can be a useful diagnostic of the bipolar winds in symbiotic stars." + }, + "0112/astro-ph0112254_arXiv.txt": { + "abstract": "We report a high-resolution (R=3000-4000) spectroscopic observation of the DA white dwarf G191-B2B in the extreme ultraviolet band $220-245$\\ \\AA . A low-density ionised He component is clearly present along the line-of-sight, which if completely interstellar implies a He ionisation fraction considerably higher than is typical of the local interstellar medium. However, some of this material may be associated with circumstellar gas, which has been detected by analysis of the CIV absorption line doublet in an HST STIS spectrum. A stellar atmosphere model assuming a uniform element distribution yields a best fit to the data which includes a significant abundance of photospheric He. The 99-percent confidence contour for the fit parameters excludes solutions in which photospheric He is absent, but this result needs to be tested using models allowing abundance gradients. ", + "introduction": " ", + "conclusions": "" + }, + "0112/astro-ph0112124_arXiv.txt": { + "abstract": "If a gravitational microlensing event is caused by a widely separated binary lens and the source approaches both lens components, the source flux is successively magnified by the individual lenses: double microlensing events. If events are observed astrometrically, double lensing events are expected to occur with an increased frequency due to the long range astrometric effect of the companion. We find that although the trajectory of the source star image centroid shifts of an astrometric double lensing event has a distorted shape from both of the elliptical ones induced by the individual single lens components, event duplication can be readily identified by the characteristic loop in the trajectory formed during the source's passage close to the companion. We determine and compare the probabilities of detecting double lensing events from both photometric and astrometric lensing observations by deriving analytic expressions for the relations between binary lensing parameters to become double lensing events. From this determination, we find that for a given set of the binary separation and the mass ratio the astrometric probability is roughly an order higher than the photometric probability. Therefore, we predict that a significant fraction of events that will be followed up by using future high precision interferometeric instruments will be identified as double lensing events. ", + "introduction": "One of the most important characteristics of microlensing light curve is that it does not repeat. However, if an event is caused by a widely separated binary lens and the source approaches both lens components, the source flux is successively magnified by the individual lenses: double microlensing events \\citep{stefano96}. Due to the special geometric condition, however, the chance to become a double lensing event is rare. Therefore, double lensing events have been neglected in the previous and current microlensing searches \\citep{udalski00, alcock00, derue01, bond01}. Although lensing events have until now been observed only photometrically, they can also be observed astrometrically by the use of high precision interferometric instruments that will be available in the near future, such as those to be mounted on space-based platforms, e.g.\\ the {\\it Space Interferometry Mission} (SIM) and the {\\it Global Astrometric Interferometer for Astrophysics} (GAIA), and those to be mounted on 10m class ground-based telescopes, e.g.\\ Keck and VLT. If an event is astrometrically observed by using these instruments, one can measure the displacement of the source star image centroid position with respect to its unlensed position (centroid shifts, $\\deltavec$). Once the trajectory of $\\deltavec$ is measured, the lens mass can be better constrained \\citep{miyamoto95, hog95, walker95, paczynski98, boden98, han99}. Recently, astrometric microlensing observation is accepted as one of the SIM long term projects (A.\\ Gould, private communication), and thus astrometric followup observations of events detected from the ground-based photometric surveys will become a routine process. One important characteristic of astrometric lensing behavior is that the astrometric effect endures to a large lens-source separation where the photometric effect is negligible \\citep{miralda96}. Then, it is expected that the chance for the source trajectory to enter the astrometrically effective lensing region of the companion will be larger, and thus the probability of detecting astrometric double lensing events\\footnote{We define the ``astrometric double lensing event'' as ``the event where the path of the source star passes through both of the astrometrically effective lensing regions of the individual lens components''.} will also be larger. In this paper, we investigate the general properties of astrometric double lensing events and estimate and compare the probabilities of detecting double lensing events from both photometric and astrometric lensing observations. The paper is organized as follows. In \\S\\ 2, we investigate the general properties of double lensing events by presenting and comparing the centroid shift trajectories and the light curves of events caused by an example wide separation binary. In \\S\\ 3, we derive analytic expressions for the relations between binary lensing parameters to become photometric and astrometric double lensing events and estimate the detection probabilities by using these relations. In \\S\\ 4, we summarize new findings and conclude. ", + "conclusions": "We have investigated the properties of double lensing events expected to be identified from future astrometric lensing observations. From this investigation, we find that although the centroid shift trajectory of an astrometric double lensing event has a distorted shape from those of the elliptical ones induced by the individual lens components, the event duplication can be readily identified from the characteristic loop formed during the source's approach close to the companion. We have also determined and compared the probabilities of detecting double lensing events from both photometric and astrometric lensing observations by deriving analytic expressions for the relations between the binary lensing parameters to become double lensing events. From this determination, we find that for a given set of the binary separation and the mass ratio, the probability to detect astrometric double lensing events is roughly an order high than the probability to detect photometric double lensing events. Therefore, we predict that a significant fraction of events that will be followed up by using future high precision interferometeric instruments will be identified as double lensing events." + }, + "0112/astro-ph0112312_arXiv.txt": { + "abstract": "We study the impact of neutrino oscillations on the supernova neutrino signal in the Large Volume Detector (LVD). The number of expected events for a galactic supernova ($D=10$ kpc) is calculated, assuming neutrino masses and mixing that explain solar and atmospheric neutrino results. The possibility to detect neutrinos in different channels makes LVD sensitive to different scenarios for $\\nu$ properties, such as normal or inverted $\\nu$ mass hierarchy, and/or adiabatic or non adiabatic MSW resonances associated to $U_{e3}.$ Of particular importance are the charged current (c.c.) reactions on $^{12}{\\rm C}$: oscillations increase by almost one order of magnitude the number of events expected from this channel. \\vskip-3.5ex ", + "introduction": "In spite of the lack of a ``standard'' model of the gravitational collapse of a massive star, some features of its dynamics and, in particular, of the correlated neutrino emission appear to be well established. At the end of its burning phase a massive star ($M \\gtrsim 8 M_{\\odot}$) explodes into a supernova (SN), originating a neutron star which cools emitting its binding energy $E_B\\sim 3\\cdot10^{53}$ erg mostly in neutrinos. The largest part of this energy, almost equipartitioned among neutrino and antineutrino species, is emitted in the cooling phase: $E_{\\bar\\nu_e} \\sim E_{\\nu_e} \\sim E_{\\nu_x} \\sim E_B/6$ ($\\nu_x$ denotes generically $\\nu_\\mu,\\bar{\\nu}_\\mu,\\nu_\\tau,\\bar{\\nu}_\\tau$). The energy spectra are approximatively Fermi-Dirac, but with different temperatures, since $\\nu_e,$ $\\bar{\\nu}_e$ and $\\nu_x$ have different couplings with the stellar matter: $T_{\\nu_e} I_A \\delta\\Omega |\\dot{\\Omega}| \\sim 10^{41} |\\dot{\\Omega}| ergs \\; s^{-1}. \\end{equation} Here $\\omega$ is the lag between the rotation rates of the crust and the pinned superfluid. $\\dot{E}_{diss}$ determines the thermal luminosities of older neutron stars. (ii) The result, from the fits to the data, that $I_A / I_c$, $I_B / I_c$ are $\\leq$ 10$^{-2}$ means the core superfluid is strongly coupled to the outer crust on short timescales, which in turn implies that precession would be strongly damped. The applicability of this simple model to all pulsar glitches and interglitch behaviour will imply that points (i) and (ii) are relevant to all neutron stars. ", + "conclusions": "" + }, + "0112/astro-ph0112130_arXiv.txt": { + "abstract": "We present a theoretical approach to infer information about RR Lyrae variables from the morphology of their light curves. The method, already successfully applied to the field first overtone variable U Comae, is now tested on a RR$_{\\mathrm ab}$ variable in the globular cluster $\\omega$~Cen. We show that, with this method, it could be possible to give an estimate of the distance modulus of the cluster with an uncertainty not larger than $\\pm$0.1 mag. The predicted variable luminosity and mass are well within the range of theoretical evolutionary expectations giving the ``feeling'' of a full compatibility between pulsational and evolutionary theory. However, before completely relying on the method, further tests are needed; the goal is that a well tested and well calibrated pulsational theory could provide reliable distance moduli from just one (or few) RR Lyrae stars. ", + "introduction": "As well known, $\\omega$~Cen is not only the largest globular cluster in our Galaxy but also contains a huge amount of RR Lyrae variables, of which more than 150 have already been detected in the cluster region. According to these features, $\\omega$~Cen could appear as an excellent target for investigating the evolutionary status of cluster stars and, in particular, to link the pulsational properties of RR Lyrae variables to their evolutionary parameters. Moreover, one can take advantage of quite a rich set of observational data as the beautiful light curves provided by Kaluzny et al. (1997) for 131 RR Lyrae stars, within the frame of the OGLE experiment. \\begin{figure}[h] \\plottwo{castellani_fig3a.eps}{castellani_fig3b.eps} \\caption{Left panel: Metallicity distribution, from Rey et al. (2000), for the Kaluzny et al. (1997) sample of RR Lyrae variables. The arrow indicates the position of the peak distribution. Right panel: Metallicity distribution, from Suntzeff \\& Kraft (1996) for giants and subgiant stars in $\\omega$~Cen.} \\label{istoFe} \\end{figure} \\begin{figure}[h] \\vspace{-2cm} \\plotfiddle{castellani_fig4.eps}{8cm}{0}{50}{50}{-150}{-50} \\vspace{1cm} \\caption{Mean visual magnitude distribution for the RR Lyrae stars in the Kaluzny et al. (1997) sample for which the metallicity estimate from Rey et al. (2000) is available. The arrows indicate the range of visual magnitude covered by the bulk or RR Lyrae variables.} \\label{NV} \\end{figure} \\begin{figure}[h] \\plottwo{castellani_fig5b.eps}{castellani_fig5a.eps} \\caption{Left panel: Zero Age Horizontal Branch (ZAHB) position for star of the labelled masses and Z=0.0004 from present computations. The central He burning evolution of selected masses is also shown. The horizontal lines indicate the range of absolute visual magnitude for the bulk of $\\omega$~Cen RR Lyrae stars (see text). The vertical dashed lines indicate the boundaries of the instability strip. Right panel: the same as the in left panel but for a metallicity Z=0.001. The ZAHB position of Z=0.0004 RR Lyrae variables is also shown for comparison. Evolutionary tracks from Cassisi et al. (1998).} \\label{ZAHB} \\end{figure} \\begin{figure}[h] \\plotone{castellani_fig6.eps} \\caption{Comparison between theory (lines) and observation (circles) for the first overtone variable U Comae, as taken from Bono, Castellani \\& Marconi (2000). From left to right the panels refer to photometric data in U,B,V (Heiser A.M. 1996) and K (Fernely, Skillen \\& Burki 1993).} \\label{striscia} \\end{figure} However, $\\omega$~Cen definitely discloses the rare peculiarity of member stars with a noticeable spread of metallicity. In probable connection with such a feature, one also finds a large spread of RR luminosities and a rather peculiar distribution of RR periods, casting some doubts on the classification of this cluster as so a ``bona fide\" Oosterhoff II. The occurrence of this metallicity spread obviously complicates things, but it offers the exciting opportunity of studying the metallicity effects on the luminosity of Horizontal Branch stars, which is a long debated question connected with relevant issues as globular cluster distances and ages. Our attention on the problem was raised by the recent paper by Rey et al. (2000), who presented metallicity estimates for 131 RR Lyrae variables in $\\omega$~Cen, as based on the use of the {\\em hk} index. Some general features of the cluster RR Lyrae population have been already discussed in that paper, as well as in a following paper by Clement \\& Rowe (2000). In this context, we approached the problem on the theoretical side, aiming to explore the degree of concordance between the quoted sets of relevant observational data and the predictions of current theories concerning both stellar evolution and stellar pulsation. However, when entering into this matter one finds that the available estimates for both the RR luminosities and metallicities, though representing an highly valued step in our knowledge of cluster stars, still leave some room for uncertainties. As a matter of the fact, by comparing field to field mean magnitudes of RR Lyrae stars, as given in Figure 1 (left panel) from the paper by Kaluzny et al. (1997), one finds that discrepancies can reach $\\approx$0.1 mag. Thus one is dealing with rather accurate light curves but with a not negligible uncertainty in the magnitude zero point. The right panel of the same Figure 1 shows that a not negligible uncertainty affects also metallicity estimates. Bearing in mind as a warning such an ``imperfect\" observational scenario, in the following we will discuss observational data to the light of current theoretical predictions. ", + "conclusions": "" + }, + "0112/astro-ph0112240_arXiv.txt": { + "abstract": "We observed the 1991 October outburst of EF Peg. Prominent superhumps with a period of 0.08705(1) d were observed, qualifying EF Peg as a being long-period SU UMa-type dwarf nova. The superhump period showed a monotonous decrease during the superoutburst, which makes a contrast to the virtually zero period change observed during the 1997 superoutburst of the same object. Large-amplitude, and highly coherent quasi-periodic oscillations (super-QPOs) were observed on October 18, when superhumps were still growing in amplitude. Most strikingly, the QPOs showed a rapid decrease of the period from 18 m to 6.8 m within the 3.2-hr observing run. Such a rapid change in the period has not been observed in any class of QPOs in cataclysmic variables. We propose a hypothesis that the rapid decrease of the QPO period reflects the rapid removal of the angular momentum from an orbiting blob in the accretion disk, via the viscosity in a turbulent disk. A brief comparison is given with the QPOs in X-ray binaries, some of which are known to show a similar rapid decrease in the periods. ", + "introduction": "\\subsection{Quasi-periodic oscillations} Quasi-periodic oscillations (QPOs) are short-period, quasi-periodic oscillations widely observed in accreting binary systems, such as cataclysmic variables (CVs) and X-ray binaries (XBs; see \\citet{vanderkli89QPOreview}; more recent reviews particularly stressing on ``kilohertz QPOs\" include \\authorcite{vanderkli99kHzQPOreview} (\\yearcite{vanderkli99kHzQPOreview}, \\yearcite{vanderkli00kHzQPOreview}). There are two major types of ``quasi-periodic\" oscillations in CVs: dwarf nova oscillations (DNOs) and (in a narrower sense) QPOs (for a review, see \\authorcite{war95book} \\yearcite{war95book}). DNOs are oscillations observed only in dwarf nova outbursts, having periods 19--29~s, and have long (several tens to $\\sim$100 wave numbers) coherence times (\\cite{rob73DNO}; \\cite{szk76DNO}; \\cite{pat81DNO}; \\cite{hil80DNO}). The amplitudes of DNOs in optical wavelengths are generally small (usually $<$0.01 mag). Several models have been proposed to account for DNOs (and QPOs in general), including vertical or radial oscillations of the accretion disk \\citep{kat78QPO}, reprocessing of the light by the orbiting blobs \\citep{pat79aeaqr}, non-radial pulsations of the accretion disk (\\cite{pap78CVQPO}; \\cite{vanhor80DNQPO}), radial oscillation of the accretion disk (\\cite{cox81QPO}; \\cite{blu84QPO}; \\cite{oku91QPO}; \\cite{oku92QPO}), excitation of trapped oscillations around the discontinuity of physical parameters \\citep{yam95DNoscillation}, and oscillation of the boundary layer of the accretion disk (\\authorcite{col98CVQPO} \\yearcite{col98CVQPO}, \\yearcite{col00CVQPO1}a,b). QPOs (other than DNOs\\footnote{We use the term ``QPOs\" in this sense in the rest of this paper.}) are more widely seen in CVs, having periods 40~s to several hundred seconds, have much shorter (usually less than $\\sim$10 wave numbers) coherence lengths. These less coherent QPOs have phenomenological similarities with QPOs in XBs, both in their coherence lengths and in the characteristic periods scaled by the Keplerian timescales at the inner edge of the accretion disk. A potentially new class of QPOs (super-QPOs) were discovered during the 1992 superoutburst of SW UMa \\citep{kat92swumasuperQPO}. These QPOs were only observed during the particular stages (usually the initial stage) of superoutbursts (for a review of SU UMa-type dwarf novae and superoutbursts, see \\authorcite{war95suuma} \\yearcite{war95suuma}), and have periods of several hundred seconds, comparable to those of QPOs other than DNOs. The most striking features of super-QPOs were the large amplitudes (up to 0.2 mag) and the long coherence, lasting at least several tens of wave numbers \\citep{kat92swumasuperQPO}. The super-QPOs observed by \\citet{kat92swumasuperQPO} consisted of sinusoidal components and short, deep dips. \\citet{kat92swumasuperQPO} also suggested that the QPOs reported by \\citet{rob87swumaQPO} showed only the sinusoidal components, but sharing all the other characteristics with super-QPOs observed during the 1992 superoutburst. \\citet{kat92swumasuperQPO} proposed that an orbiting blob, which incidentally eclipsed the central part of the accretion disk, was responsible for the large-amplitude super-QPOs. \\citet{kat92swumasuperQPO} suggested that the reason why super-QPOs are only observed during the particular stage of superoutburst may be that such a blob can be formed under the enhanced dissipation (i.e. turbulent condition) caused by the tidal instability, which is responsible for superhumps (\\cite{whi88tidal}; \\cite{hir90SHexcess}). Studies of super-QPOs have been, however, hindered by the paucity of examples. EF Peg is the only other example, which showed the remarkable time evolution of super-QPOs during the 1991 superoutburst. \\subsection{EF Peg} EF Peg was originally discovered as a variable star \\citep{hof35an255407}, who suggested the Mira-type classification. \\citet{esc37efpeg} recorded two additional (bright and faint ones) maxima, and derived a period period of 157 d in combination with the maxima recorded by \\citet{hof35an255407}. However, \\citet{san50efpeg} only recorded an apparently nonvariable star of 13.5--14.0 photographic magnitude on 157 plates. \\citet{tse79efpeg} finally revealed that the true EF Peg is a normally fainter companion to a nonvariable star (labeled as $w$ in \\cite{tse79efpeg} and \\cite{how93efpeg}). \\citet{tse79efpeg} suggested that EF Peg is a dwarf nova with a large outburst amplitude. Their observations showed the existence of two types of --- short and long --- outbursts, which strongly suggested an SU UMa-type dwarf nova (for a recent review of SU UMa-type stars and their observational properties, see \\cite{war95suuma}). Upon this information, the Variable Star Observers League in Japan (VSOLJ) called for a monitoring campaign for an outburst since 1985. \\citet{ges88efpeg} searched Sonneberg plates, and found nine outbursts between 1928 and 1986, confirming the classification by \\citet{tse79efpeg}. No firm outbursts above a magnitude of 12 had been recorded until the discovery of an outburst by Patrick Schmeer on 1991 October 15.779 UT at visual magnitude 10.9 \\citep{sch91efpegiauc}. Since then, only two secure outbursts have been observed in 1995 January and 1997 October, according to the reports to the VSNET Collaboration\\footnote{ $\\langle$http://www.kusastro.kyoto-u.ac.jp/vsnet/$\\rangle$ }, and no further outburst has been observed up to 2001 July. These observations indicate the intrinsically low outburst frequency of this object. \\citet{pol98TOAD} even suggested that EF Peg is past the period minimum of the CV evolution. The first detection of superhumps was made by \\citet{kat91efpegiauc} on 1991 October 18. More information of this superoutburst is given in \\citet{how93efpeg}. ", + "conclusions": "\\subsection{Superhump period change} As described in Section \\ref{sec:pdot}, EF Peg showed a decrease of the superhump period at a rate of $\\dot{P}$ = -4.4$\\pm$0.6 $\\times$ 10$^{-6}$ d cycle$^{-1}$, or $\\dot{P}/P$= -5.1(0.7) $\\times$ 10$^{-5}$. This negative period derivative has a typical value for long-period SU UMa-type dwarf novae (cf. \\cite{kat01hvvir}; \\cite{pat93vyaqr}; \\cite{kat98super}). However, the same EF Peg showed virtually zero period change during the 1997 superoutburst (Matsumoto et al., in preparation). While only a few SU UMa-type systems have been observed during different superoutbursts, all known systems showed relatively constant $\\dot{P}$ even in different superoutbursts \\citep{kat01hvvir}. If such a large variation of $\\dot{P}$ in EF Peg is confirmed by future observations, EF Peg would be a unique system with regards to the stability of $\\dot{P}$. \\citet{kat98super} proposed that positive $\\dot{P}$ in some short-period SU UMa-type systems may be a result of outward propagation of the eccentricity wave (cf. \\cite{lub92SH}; \\cite{bab00v1028cyg}). The large variation of $\\dot{P}$, proposed by the present observation of EF Peg, may reflect the different locations where tidal instability ignites. \\subsection{Super-QPOs} As described in Section \\ref{sec:superqpo}, EF Peg showed remarkable quasi-periodic oscillations (QPOs) during the earliest stage of the superoutburst, i.e. the growing stage of superhumps (October 18). These QPOs had disappeared when the amplitude of superhumps reached the maximum (October 20). As seen in \\ref{fig:qpooc}, the maximum jump in the QPO phase was less than 0.002 d, indicating that the oscillations were essentially coherent during the entire observing run covering 20 QPO cycles, although a strong global variation of the period was present. Such a long coherence of oscillations is a signature of super-QPOs \\citep{kat92swumasuperQPO}. Although the amplitudes of the present QPOs were smaller than that recorded in SW UMa during the 1992 superoutburst \\citep{kat92swumasuperQPO}, they were larger than those of QPOs observed by \\citet{rob87swumaQPO} for the same star. The observed QPO profile and amplitude (figure \\ref{fig:qpoprof}) resemble those of the sinusoidal component of \\citep{kat92swumasuperQPO}. From the transient appearance in the early stage of a superoutburst and the large amplitudes, we consider the present QPOs as a variety of super-QPOs. The same explanation in \\citet{kat92swumasuperQPO} could apply to the present large-amplitude QPOs, except that the present QPOs incidentally lack the dip component. The rapid decrease of the QPO period (figure \\ref{fig:qpovar}, \\ref{fig:qpooc}) is the unique characteristic of the present QPOs. Such a large variation of the QPO frequency within a short time (less than 3.2 hr) has never been observed in other QPOs of CVs. Furthermore, the period decrease occurred continuously (Section \\ref{sec:superqpo}), rather than a sudden mode switch between different periods. The phenomenon is hard to explain, if these QPOs arise from oscillations of the accretion disk, or from the beat between the rotation of the magnetic white dwarf and the slowly rotating part of the accretion disk. Such a rapid decrease of the QPO period can be more naturally understood if the rotating blob (cf. \\cite{kat92swumasuperQPO}), which rapidly lost its angular momentum via the viscosity in a turbulent disk. If we assume a 1 M$_\\odot$ primary and nearly Keplerian orbits, a decrease of the period from 18 m to 6.8 m corresponds to a decrease of the radius of the orbit from 1.6 10$^{10}$ to 8.3 10$^{9}$ cm. This decrease in 3.2 hr corresponds to a radial velocity of 6.7 10$^{5}$ cm s$^{-1}$ (1/140 -- 1/200 of the orbiting velocity), which seems to be a reasonable value for accretion during outburst. \\subsection{Comparison with QPOs in XBs} Although such a rapid decrease of the QPO period is unique among CVs, there are a few known similar examples in oscillations (or large-amplitude QPOs) in some XBs. The best-known example is seen in the Rapid Burster (MXB 1730-335). The first indication of such oscillations during its type II bursts from the Rapid Burster was found with the GINGA satellite \\citep{taw82RapidBurster}. \\citet{dot90RapidBursterQPO} systematically studied these oscillations, and found the evidence of strong period changes. The best example of the ``naked-eye\" visibility of large-amplitude oscillations and the rapid decrease of their periods can be found in \\authorcite{lub92RapidBursteroscillation} (\\yearcite{lub92RapidBursteroscillation}) [see also \\authorcite{lub91RapidBursterQPO} \\yearcite{lub91RapidBursterQPO},\\yearcite{lub92RapidBurstertype2burst}; for a recent observational review, see also \\cite{gue99RapidBurster}). The second example of similar oscillations was found in the bursting X-ray pulsar GRO J1744-28 (\\cite{zha96j1744QPO}; \\cite{lew96RapidBursterj1744}). The characteristics of the oscillations in GRO J1744-28 was studied in detail by \\citet{kom97j1744QPO}. Although the nature of such oscillations are still poorly understood, several attempts have been made to understand these two unique objects. In the case of GRO J1744-28, \\citet{can96j1744} considered an interplay between radial and vertical energy transport in the accretion disk, causing oscillations in the fraction of the gas pressure to the total pressure. \\citet{can96j1744} showed that these oscillations eventually lead to a Lightman-Eardley instability \\citep{LightmanEardleyInstability}, which \\citet{can96j1744} considered to be the cause of Type II bursts. \\citet{can97j1744} further showed that the characteristic QPO frequency can be reproduced by a selection of parameters. \\citet{abr95XBQPO} suggested that low-frequency QPOs in some XBs (including those of the Rapid Burster) and black-hole candidates can be produced via dissipation of energy in the coronal region above the accretion disk. While these models were able to reproduce some parts of characteristics of QPOs (and bursts) observed in these systems, they have not yet succeeded in explaining a rapid decrease of QPO periods; the cause of this phenomenon is thus still an open question. These models for the Rapid Burster and GRO J1744-28 assume the strong radiation pressure near the inner edge of the accretion disk. Since such a condition is hard to achieve in a CV disk, the mechanisms proposed for QPOs in XBs would not be directly applicable to the present QPOs in EF Peg. However, the existence of phenomenologically similar QPOs both in CVs and XBs may provide future insight for the fundamental understanding the underlying mechanisms." + }, + "0112/astro-ph0112289_arXiv.txt": { + "abstract": "On the basis of theoretical evolutionary expectations, we develop a Galactic model reproducing star counts and synthetic color-magnitude diagrams of field stars, which include the white dwarf (WD) population. In this way we are able to evaluate the expected occurrence of WDs in deep observations at the various photometric bands and Galactic coordinates, discussing the contribution of the WDs of the various Galactic components. The effects on the theoretical predictions of different WD evolutionary models, ages, initial mass functions and relations between progenitor mass and WD mass are discussed. ", + "introduction": "Our Galactic model is a three components (spheroid, disc, thick disc) code which closely follows the ``classic'' Galactic models by Bahcall \\& Soneira (1984, B\\&S) and Gilmore \\& Reid (1983, see also Paper I and references therein) concerning the spatial density distribution of stars. However, it relies on suitable assumptions on the evolutionary status and on the initial mass function (IMF) of the various Galactic populations to reproduce the luminosity functions used as an (observational) input in Paper I. We recall here the main characteristics of the model referring, for a detailed discussion, to Castellani et al. (2001, Paper II). Predicted results are obtained by randomly generating star masses according to the adopted IMF (Kroupa 2001) and by using stellar models to derive luminosities in the selected bands for each given value of the stellar mass and age. Spheroid stars are assumed to be almost coeval and thus they are reproduced by populating a suitable theoretical isochrone (age= 12 Gyr, Z= 0.0002, Y= 0.23), while for both the thick disc and the disc one has to take into account prolonged episodes of star formation. Thus for these two components star masses and ages are both randomly generated, the mass distribution reproducing the selected IMF, while a flat age distribution is adopted within the range assumed for each component. For the disc we chose a constant star formation rate (SFR) from 50 Myr to 9 Gyr to populate evolutionary tracks (Z=0.02 Y=0.27) until the asymptotic giant branch and the WD cooling sequences. For the thick disc, following Gilmore, Wyse, \\& Jones (1995) and Norris (1999), we assume a metallicity of $\\sim$Z=0.006 and a SFR centered at $\\sim$ 10 Gyr with a spread of few Gyr. Thanks to this theoretical approach the model spontaneously predicts the occurrence of stars in the various evolutionary phases. ", + "conclusions": "We described our galactic model able to reproduce star counts and synthetic CMDs of field stars from the main sequence to the WD evolutionary phase at the various photometric bands and Galactic coordinates. The main goal is the introduction of disc/thick disc and spheroid WD population in an evolutionary consistent way. We found that the WD population is barely sensitive to a change of theoretical WD models or to a variation of the adopted relation between the WD mass and the progenitor mass." + }, + "0112/astro-ph0112076_arXiv.txt": { + "abstract": "I present $N$-body simulations of isolated and interacting galaxies, made on GRAPE machines. In particular I discuss the formation and evolution of $N$-body bars and compare their properties to those of bars in early-type and late-type galactic discs. I argue that the halo can help the bar grow, contrary to previous beliefs, by taking positive angular momentum from it via its resonant stars. I then focus on the interaction and subsequent merging of a barred disc galaxy with a spheroidal satellite. The evolution depends strongly on the mass (density) of the satellite and may lead to its destruction or to the destruction of the bar. ", + "introduction": "GRAPE machines (Makino, these proceedings) have given a substantial boost to the study of galactic dynamics by making state of the art $N$-body simulations possible at a relatively low cost and by allowing a simulation environment which is much more flexible than that of supercomputing centers. Machines with even numbers, e.g. GRAPE-4 (Makino et al. 1997) and GRAPE-6 (Makino, these proceedings), have a very high accuracy arithmetic and calculate also the derivative of the force, thus allowing a more accurate integration scheme to follow the evolution, as well as calculations with little or no softening. In the field of galactic dynamics they are particularly suited for studying the nuclei of galaxies, the effect of black holes, as well as for problems where softening should be limited or avoided, or for other cases where high precision is necessary. Machines with odd numbers, as GRAPE-3 (Ebisuzaki et al. 1993) and GRAPE-5 (Kawai et al. 2000), have a more limited accuracy and do not provide the derivatives of the force. Nevertheless the accuracy of GRAPE-3 was shown to be amply sufficient for a large number of problems on galactic dynamics (e.g. Athanassoula et al. 1998) and GRAPE-5 is even more accurate (Kawai et al. 2000). They are thus very well suited for studies of the evolution of isolated and interacting galaxies, as well as of galaxy systems. They have therefore been widely used in these topics, as witnessed by the ever-increasing number of publications based on GRAPE simulations. The field of research thus covered is too broad to be usefully covered by a single review, so I will limit myself here to a few, largely unpublished, results obtained recently with the Marseille GRAPE systems. In sections~2 and 3 I will present results on the formation of bars and on the secular evolution of barred disc galaxies. In section~4 I will discuss the interaction and subsequent merging of a barred disc galaxy and a small elliptical companion. A few concluding remarks are given in section~5. ", + "conclusions": "\\label{sec:conclusions} In the above I presented results of simulations made with the Marseille GRAPE systems. I started with simulations of the formation and evolution of a bar in an isolated disc galaxy. These show that strong bars can form in systems with a sizeable halo component within the disc region. I find very good agreement between the properties of the $N$-body bars and those of bars in galaxies. In particular I presented arguments that link $N$-body bars that grew in a environment with a considerable halo component with early type bars, and bars growing in a disc-dominated environment with late type bars. Such simulations stress the role of the halo in the development of the bar instability and show that the halo can stimulate bar formation, contrary to the common belief that it will always quench it. The reason is that the halo responds to the bar and a considerable fraction of its particles are in resonance with it. This of course would have been missed by simulations treating the halo as a rigid component. I then discussed the interaction and subsequent merging of a barred disc galaxy with its spherical satellite. If the satellite has high mass (density), then it will spiral rapidly towards the center of the disc, destroy the bar and occupy the center of the target, thus forming a bulge, or contributing to one. Such interactions will cause evolution along the Hubble sequence, from late to early type discs. During this process the disc thickens considerably, but also expands, so that its final shape is still that of a disc, albeit somewhat thicker than the initial one. If the companion has low mass (density), then it will get disrupted as it spirals inwards towards the center of the target. If the orbit of the companion is initially at an angle with the plane of the disc, then the disc tilts but is not destroyed. In the case of the high mass (density) companion, the tilt angle is not far from that of the plane of the initial orbit of the companion. The above results, and many others not shown here, argue that GRAPE is an ideal tool for simulations of galaxies and galaxy systems. In all the above simulations the number of particles was of the order of a million, or more than that. This is now a standard number for most of our GRAPE simulations. GRAPE can of course handle considerably larger numbers, provided its front end has a sufficient memory, and we have used such numbers whenever necessary. GRAPE boards permit fast and accurate calculations of the gravitational potential and forces, and are not limited by the specific geometry of the particle distribution. Their price is a very small fraction of that of supercomputers. Both high precision direct summation codes and tree codes have been developed for them. They offer a flexible simulation environment. They are thus ideally suited for small research groups ready to invest themselves in the GRAPE approach. The close collaboration between the members of the international community of GRAPE users contributes further to the success of these machines." + }, + "0112/hep-ph0112069_arXiv.txt": { + "abstract": "{We propose in this paper a scenario of spontaneous baryogenesis in cosmological models of Quintessence by introducing a derivative coupling of the Quintessence scalar $Q$ to the baryon current $J_B^{\\mu}$ or the current of the baryon number minus lepton number $J_{B-L}^{\\mu}$. We find that with a dimension-5 operator ${\\partial_\\mu Q} J_{B-L}^{\\mu}$ suppressed by the Planck mass $M_{pl}$ or the Grand Unification Scale $M_{GUT}$, baryon number asymmetry $n_B/s \\sim 10^{-10}$ can be naturally explained {\\it via} leptogenesis. We study also the isocurvature baryon number fluctuation generated in our model.} ", + "introduction": " ", + "conclusions": "" + }, + "0112/hep-ph0112319_arXiv.txt": { + "abstract": "We estimate in this paper the errors in the zenith angle distribution for the charged current events of the solar neutrinos caused by the uncertainty of the earth electron density. In the model of PREM with a $5\\%$ uncertainty in the earth electron density we numerically calculate the corrections to the correlation between $ [N]_5 /[N]_2 $ and $ [N]_2 /[N]_3 $, and find the errors notable. ", + "introduction": " ", + "conclusions": "" + }, + "0112/astro-ph0112526_arXiv.txt": { + "abstract": "We explore various pitfalls and challenges in determining the equation-of-state ($w$) of dark energy component that dominates the universe and causes the current accelerated expansion. We demonstrated in an earlier paper the existence of a degeneracy that makes it impossible to resolve well the value of $w$ or its time-derivative with supernovae data. Here we consider standard practices, such as assuming priors that $w$ is constant or greater than $-1$, and show that they also can lead to gross errors in estimating the true equation-of-state. We further consider combining measurements of the cosmic microwave background anisotropy and the Alcock-Paczynski test with supernovae data and find that the improvement in resolving the time-derivative of $w$ is marginal, although the combination can constrain its present value perhaps to 20 percent uncertainty. ", + "introduction": "Measurements of Type IA supernovae have shown that the expansion of the universe is accelerating \\cite{SCP,HZS}, suggesting that most of the energy density of the universe consists of some form of dark energy with negative pressure \\cite{EOS}. Combining measurements of the cosmic microwave background anisotropy and observations of large-scale structure provides important corroborating evidence.\\cite{OS,Bahcall} Two candidates for the dark energy are a cosmological constant (or vacuum density) and quintessence,\\cite{CDS} a time-varying, spatially inhomogeneous component. In a previous paper \\cite{Maor:2001jy} (Paper I), we addressed the question of whether supernova measurements can be used to measure the equation of state (EOS) of the negative pressure component, the ratio $w$ of the pressure to the energy density. The issue is important because $w=-1$ for a cosmological constant whereas $w$ takes on different values and can be significantly time-varying in the case of quintessence.\\cite{CDS,other} Under the assumption that $w$ is constant, its value can be determined to better than 5 per cent by measuring several thousand supernovae distributed equally between red shift $z=0$ and $z=2$. However, we showed that a degeneracy opens up if $w$ is time-dependent which makes it impossible to determine accurately the current value of $w$ or its time-derivative. The cause of the degeneracy is that supernovae measure luminosity distance, which is related by a multi-integral expression to the EOS as a function of red shift, $w(z)$. Widely different $w(z)$ can have the same multi-integral value. The purpose of this paper is to explore some pitfalls and challenges in determining the EOS and its time variation using supernovae. For example, we shall show how the standard practice of considering only models with $w\\ge -1$ or only models with constant $w$ when doing likelihood analyses can lead to grossly incorrect results. For example, we will illustrate cases where the standard practice will suggest that $w$ is near -1 or much more negative than $-1$ when, in fact, $w$ is significantly greater than -1 and rapidly time varying. We shall show that a non-zero value of $dw/dz$ is more easily detected if $dw/dz>0$ than if $dw /dz <0$. We shall also contrast measuring $w$ for the negative pressure component alone ($w_Q$) versus the mean value for the total energy density (including ordinary and dark matter), $w_T$. The degeneracy problem is less severe for $w_T$, but this parameter provides less useful information. We consider possibilities of breaking the degeneracy between $w$ and $d w/dz$ by combining supernovae results with either cosmic microwave background anisotropy measurements and/or the Alcock-Paczynski test. We shall show that neither additional test significantly improves the measurement of the time-variation of $w$, although optimistic assumptions about the Alcock-Paczynski test suggest that the current value of $w$ can be measured to within 20 percent or so. We conclude that a new, yet to be found test has to be devised to resolve well the cosmic EOS and its time variation. We stress that with current data it is possible to determine the EOS to about a factor of two. For a future experiment to significantly enhance the determination of the EOS, and enable the distinction between a constant EOS and a time-dependent one, it needs to resolve the equation of state at the 10\\% level or better. The results of our analysis agree with many other analyses \\cite{Podariu,astier,Weller,Chiba,Barger,Chevallier,Goliath,trentham,gudmun,weller2,Ng,Kujat}, although not always with their interpretation, and can be used to explain why some other analyses seem to indicate a superior resolving power of SN measurements alone \\cite{Hut,Staro,Boisseau,Wang}, or in combination with other measurements \\cite{Huterer,Tegmark,Corasaniti,Lovelace}. Some of the latter analyses implicitly assume unrealistic accuracy in independent determination of cosmological parameters, by not including self-consistently the uncertainty in $w(z)$ in all measurements. For example, assume a reported resolution of matter energy density which was based on assuming $w_Q=-1$. Type IA supernovae have intrinsic variability of about 0.15 in absolute magnitude, but currently the errors in measuring distant SN are above this. There are ongoing programs to extend the search to deeper redshifts and improve measurement quality. The proposed SNAP satellite plans to measure 2000 SN per year, mostly in the range of redshifts $0.1-1$ and rapidly increasing. (Sec. IV) \\item Time-variation of $w_Q$ is more easily detected if $w_Q(z)$ is an increasing function of $z$ rather than decreasing. (Sec. V) \\item To resolve $w_Q(z)$ with supernova data, an additional test is needed. Given optimistic estimates of experimental uncertainties, the Alcock-Paczynski test combined with the supernovae measurements can constraint the current value of $w$ to within 20 percent or so. However, Neither the Alcock-Paczynski test nor microwave background anisotropy measurements provide the needed resolution to constrain the time-variation. (Sec. VI) \\end{itemize} Our principal conclusion is that a new test is required to achieve the goal of measuring $w_Q(z)$. In devising a new test, the two considerations must be precision and model dependence. Thus far, among the measurements that we have considered, the measurements which are precise give constraints on $w_Q$ that are highly model dependent, leading to degeneracy problems. Tests which are not model dependent turn out to be difficult to measure precisely. So, there lies the challenge. In considering alternatives, it is critical to include practical estimates of their uncertainties. Furthermore, one must consider how the new tests, themselves, depend on $w_Q(z)$. For example, claims have been made that $\\Omega_m$ and $\\Omega_Q$ have been or will be measured very accurately by measurements of the cosmic microwave background \\cite{Turner:2001mw}. However, those estimates are based on assuming that $w_Q=-1$. Making no prior assumption about $w_Q(z)$, a degeneracy problem once again arises\\cite{degen} that spoils the resolution of $\\Omega_Q$ and $w_Q(z)$, as discussed in Sec. VI. While trying to devise a new test to determine $w_Q$, it is worth mentioning that a precise measurement of $H'$ will be extremely useful. The dependence of $w_Q$ on $H$, $H'$ (prime denotes a derivative with respect to $x=1+z$) and $\\Omega_m$ is given by $ w_Q=\\frac{\\frac{2}{3} x HH'-H^2} {H^2-\\Omega_m H_0^2\\ x^3}.$ A good measurement of $H'$ is clearly crucial to the resolution of $w_Q$, but current tests do not probe $H'$ directly. The next best option is to measure $H(z)$, and then estimate $H'$ by calculating its derivative. Obviously, this worsens the resolution for $H'$ and increases the uncertainty in $w_Q$. Three additional approaches that we have not tried yet are measuring the time dependence of structure growth on $w(z)$;\\cite{Weller3} gravitational lensing; and direct measurements of $dz/dt$ (to be discussed elsewhere\\cite{Loeb})." + }, + "0112/astro-ph0112460_arXiv.txt": { + "abstract": "A thorough analysis of well studied giant HII regions on galactic discs for which we know the ionizing stellar population, the gas metallicity and the Wolf-Rayet population, leads to photoionization models which can only match all observed line intensity ratios ([OIII] , [OII] , [NII], [SII] and [SIII] with respect to the intensity of H$\\beta$), as well as the H$\\beta$ luminosity and equivalent width if one allows for an important escape of energetic ionizing radiation. For the three regions presented here, the fractions of escaping Lyman continuum photons amount to 10 to 73 \\% and, in all cases, the larger fraction of escaping photons has energies between 13.6 and 24.4 eV. These escaping photons clearly must have an important impact as a source of ionization of the diffuse ionized gas (DIG) found surrounding many galaxies, as well as of the intergalactic medium (IGM). ", + "introduction": "The issue of whether or not ionizing star clusters are the sources of radiation responsible for the diffuse ionized gas (DIG) and, moreover, for the background radiation at low redshift (Heckman et al 2001), is still controversial. UV observations of four nearby starburst galaxies (Leitherer et al. 1995; Hurwitz et al. 1997) showed that more than 10\\% of the ionizing photons could escape from these galaxies. On the other hand, Dove et al. (2000) obtained from theoretical arguments that 7\\% of photons produced by OB associations would have to escape into the DIG to sustain it and could also be consistent with the estimated flux required to photoionize the Magellanic Stream (Weiner \\& Williams 1996; Bland-Hawthorn \\& Maloney 1999). All these aspects at low redshift are undoubtedly important to establish whether massive stars in starburst galaxies can rival QSO's as sources of the metagalactic background radiation at high redshifts (z $>$ 3; see Ricotti \\& Shull 2000). Furthermore, an estimate of the escape fraction at low redshift is to provide an upper limit to the expected escape fraction at high redshift, when the universe was much denser. Beckman et al. (2000) presented evidence, based on the analysis of the H$\\alpha$ luminosity function of HII region populations in nearby spiral galaxies, that luminous HII regions may be matter bounded. This would imply that these regions are an important source of photons to the DIG. Here we report strong evidence that points at Giant Extragalactic HII Regions (GEHRs) to be matter bounded, and thus as an important source of ionization at large distances from their exciting stars. Our conclusions are based on evolutionary synthesis models (Leitherer et al. 1999; hereafter ST99) applied to previously analyzed GEHRs with observed Wolf-Rayet (WR) features (D{\\'\\i}az et al. 2000; Castellanos et al. 2002; hereafter referred to as Paper I). The observed parameters inside these regions, i.e. the fluxes and equivalent widths of the WR features, the emission line intensities of [OII]($\\lambda$3727), [OIII]($\\lambda$5007), [NII]($\\lambda$6584), [SII]($\\lambda$6717), and [SIII]($\\lambda$9069) relative to H$\\beta$, the observed H$\\beta$ luminosity and equivalent width, cannot be fitted simultaneously unless an important fraction of the ionizing radiation escapes from these regions. ", + "conclusions": "The observed emission lines of GEHR have been widely used to derive the properties of the ionizing radiation and link them to those of the stellar populations generating it, but always under the assumption that the ionized region is ionization bounded. In the absence of an independent constraint on the age of this stellar population, such as the presence of WR stars, successful fits to the emission line spectra of most GEHR have been found within a narrow range of ages, between 2 and 2.5 Myr (Bresolin et al. 1999; Dopita et al. 2000). In fact, if similar methods of analysis were applied to the regions modelled here, we would obtain a similar result. For these regions, however, we have accurate determinations of the gas metallicity and their age, provided by the analysis of the WR spectral features, and therefore we can synthesize the corresponding SEDs. However, the derived SEDs are unable to reproduce the observed emission line spectra unless the GEHRs are density bounded. In such a case, we have shown here that a satisfactory fit can be found in the three analyzed cases. Our results suggest that the leaking of ionizing photons do depend on the evolutionary stage of the regions. Regions 74C and CDT3 show the highest escape fractions, consistent with the prominent observed WR features in the former spectrum, and the high excitation lines ([FeII],, [FeIII]) in the latter one. However, region H13, which also shows a prominent WR feature, shows a far lower escape fraction which can be understood from its lower metallicity (0.2 solar {\\it versus} 0.5 solar in 74C and CDT3), affecting the WR phase strength. It is plausible that large photon escaping fractions might be expected in those regions that experience a hard WR phase, this one depending on both the age and metallicity. An important implication from the escape of ionizing photons concerns the EW(H$\\beta$) - age relation, that does not hold for matter bounded regions. It is a well known fact that few HII regions show EW(H$\\beta$) as large as those predicted from evolutionary synthesis models (e.g. Melnick, Terlevich \\& Eggleton, 1985; Mas-Hesse \\& Kunth 1998). The leak of Lyman continuum radiation provides a natural explanation in the correct direction. A close analysis of the energy distribution of the escaping radiation (see Table 1) shows that the largest fraction has energies between 13.6 and 24.4 eV (up to 80\\% in the case of 74C). However, if escaping to incident fractions are considered, helium ionizing photons are clearly dominant (e.g. up to 84\\% of the incident photons with energies between 24.4 and 54.4 eV escape from region 74C). This would explain why the observed emission line spectrum is so similar to that produced by a star cluster without WR stars. Furthermore, the derived large amounts of escaping ionizing photons indicate that GEHRs and stellar OB associations may significantly contribute to the ionization of the diffuse gas layers observed above the disks of spiral galaxies. Though, to our knowledge, no measurements of the diffuse interstellar gas (DIG) are available for NGC 628, NGC 1232 and NGC 4258, recent studies by Ferguson et al. (1996), Oey \\& Kennicutt (1997) and Zurita et al. (2000) stress the fact that the integrated escaping flux from disk galaxies can attain up to 10$^{41}$ erg s$^{-1}$, enough to account for the diffuse H$\\alpha$ flux and fully compatible with our derived values for single GEHRs. In our models the lower energy photons emitted by the central star cluster are more effectively absorbed by the HII region gas while a large fraction of the photons with energy between 1.8 to 4 Ry escape the nebula. Therefore, the resulting spectral energy distribution of the escaping photons is slightly harder than that of the ionizing star cluster. For the three studied regions the Q(He)/Q(H) ratio of the escaping photons is: 0.22, 0.20 and 0.19 which should be compared to 0.10, 0.14 and 0.15 for the respective ionizing clusters. One still unsolved problem for the interpretation of observations of the DIG above the galactic disk is the apparently low value of the HeI$\\lambda$5876 ${\\AA}$ /H$\\beta$ (Reynolds \\& Tufte 1995) which would be in contradiction with our derived Q(He)/Q(H) ratios. Our studied regions, however, are not representative of galactic HII regions ionized by a single star, but of high luminosity Giant Extragalactic HII Regions on the disks of spiral galaxies like, for example, NGC891. For the DIG in this galaxy the average value of HeI$\\lambda$5876 ${\\AA}$ /H$\\beta$ is 0.1 (Rand 1997). According to Bresolin, Kennicutt \\& Garnett (1999) the value of this ratio for stellar effective temperatures between 40000 and 50000 K is highly dependent of metallicity and reaches the value of 0.1 for a metallicity of about 1/4 solar. In fact, most HII galaxies show values of HeI$\\lambda$5876 ${\\AA}$ /H$\\beta$ around 0.10 (Izotov \\& Thuan 1998 and references therein). In summary, by applying evolutionary synthesis models to well-studied GEHRs with observed WR features, we find that these regions need to be matter bounded in order to reconcile model predictions with observations. The amount of escaping ionizing photons per unit time from these regions range from 49.80 $<$ log Q$_{esc}$(H) $<$ 51.51, which implies an escape fraction between 10\\% and 73\\% of the available incident photons. These fractions seem to increase with both the strength of the WR phase and the metallicity of the ISM. The implication is thus that the mass of the ionizing clusters is larger than the value derived in a straight manner from the observed H$\\alpha$ luminosity. It is also clear that the ISM is highly non-uniform and that the matter swept up in the shells produced by the mechanical energy deposited by the WR sources and other massive stars is not sufficient to trap the ionization fronts. This is also the result obtained from numerical calculations that consider both the mechanical energy of massive stellar clusters as well as their ionizing luminosity. Tenorio-Tagle et al. (1999) have shown how once the shell of swept up matter becomes Rayleigh - Taylor unstable and fragments, as it evolves out of a galaxy disk into the halo (the blowout phenomenon), it then allows not only for the venting of the hot (wind and SN) matter into the halo, but also for the leakage of a large fraction of the ionizing radiation. The latter soon establishes a giant conical HII region in the low density halo. The low densities in the halo lead to a long recombination time and thus to a large leakage of photons into the IGM. The situation changes later, once the shock is able to sweep enough halo matter, enhancing locally the number of recombinations in the expanding shell. This leads eventually to the trapping of the ionization front. Consecuently, depending on the stage of the evolution, a large fraction of the ionizing flux escapes the nebulae and is freely available to impact upon the gas at large distances from the host galaxy plane (Collins \\& Rand 2001) and it is also likely to escape the galaxy and cause an important ionization of the IGM. Clearly, further analysis of this kind must be done to infer whether or not our derived escape fractions are typical of other GEHRs." + }, + "0112/astro-ph0112183_arXiv.txt": { + "abstract": "We are conducting a large survey of star formation within $\\sim400$ nearby spiral and irregular galaxies through imaging of the H$\\alpha$ emission. We present here some of our first results from 104 of the galaxies in our sample, investigating the variation of SFR and H$\\alpha$ EWs along the Hubble sequence for these galaxies. We find a strong dependence of SFR on Hubble type, peaking for Sbc galaxies, but with large dispersions within each type. There is a possible dependence of SFR on bar presence, but none on group membership. There is an increase in EWs for later Hubble types, but with large dispersions within each type. We find no dependence of EW on bar presence, group membership or on absolute magnitude. ", + "introduction": "The H$\\alpha$ Galaxy Survey (H$\\alpha$GS) is a major survey of star formation in the local Universe, using the 1m Jacobus Kapteyn Telescope (JKT) at La Palma, and redshifted H$\\alpha$ filters to determine the quantity and spatial distribution of ionized gas in $\\sim400$ spiral and irregular galaxies within 30 h$^{-1}$ Mpc. We are using broad- and intermediate-band R filters for continuum subtraction. Our sample was selected, using the Uppsala Galaxy Catalogue (Nilson 1973) as the parent catalogue, within 5 redshift shells, but with the angular area of the selection region decreasing such that all shells have equal volumes. The well-defined selection criteria, and the large total number of galaxies, mean that it will be possible to combine the data for all shells such that galaxies with a wide range of luminosities, diameters and surface brightness are represented. Previous studies, e.g. those of Kennicutt \\& Kent (1983) and Young et al. (1996), have been weighted heavily towards large, luminous and relatively rare Sa-Sc spirals. The low luminosity, dwarf galaxies in our sample are the most populous in the Universe, and our survey will determine whether they are in fact responsible of a significant fraction of the total star formation locally. We have been awarded 14 weeks of observing time over 2000-2002 (long-term status) on the JKT to complete the H$\\alpha$ observations. We are aiming to address such questions as: \\\\- What is the total star formation rate (SFR) in the local Universe? \\\\- Which galaxy types and morphologies dominate? \\\\- Does nuclear activity (star formation or AGN) depend on galaxy properties and environment? \\\\- What is the star formation distribution within galaxies? \\\\- How does environment affect star formation rate? In this paper we will present preliminary results from SFRs and equivalent widths (EWs) so far calculated for 104 of the galaxies in our sample. These galaxies were all observed under photometric conditions. \\begin{figure} \\plotfiddle{shane_fig1.ps}{5.5cm}{270}{33}{33}{-120}{180} \\caption{Comparison of H$\\alpha$ to IRAS SFRs} \\end{figure} ", + "conclusions": "From a preliminary subset of 104 galaxies ($\\sim\\onequarter$ of the final sample) we find: \\\\- a correlation between our H$\\alpha$GS SFRs and IRAS SFRs; \\\\- a strong dependence of SFR on Hubble type, but with large dispersions within each type; \\\\- a possible dependence of SFR on bar presence for late type galaxies, but none on group membership; \\\\- an increase in EWs for later types, but large dispersions within each type; \\\\- no dependence of EW on bar presence or group membership; \\\\- no correlation between EW and absolute magnitude." + }, + "0112/astro-ph0112111_arXiv.txt": { + "abstract": "In this paper we investigate the level of agreement between observations and ``new\" Asymptotic Giant Branch (AGB) models, as produced by updating the physical inputs adopted in previous stellar computations. One finds that the new physics increases the predicted luminosity of Horizontal Branch (HB) and AGB stellar structures by a similar amount, keeping unchanged the predictions about the difference in luminosity between these two evolutionary phases. The best fit of selected globular clusters appears rather satisfactory, disclosing the relevance of the assumption on the mass of the Red Giant Branch (RGB) progenitor in assessing the distance modulus of moderately metal rich clusters. The still existing uncertainties related either to the input physics or to the efficiency of some macroscopic mechanisms, like convection or microscopic diffusion, are critically discussed, ruling out the occurrence of the so called ``breathing pulses\" during the central He exhaustion, in agreement with earlier suggestions. ", + "introduction": "The capability of current stellar models to account for all the evolutionary phases observed in stellar clusters is undoubtedly an exciting achievement which crowns with success the development of stellar evolutionary theories as pursued all along the second half of the last century. Following such a success, one is often tempted to use evolutionary results in an uncritical way, i.e., taking these results at their face values without allowing for theoretical uncertainties. However, theoretical uncertainties do exist, as it is clearly shown by the not negligible differences still existing among evolutionary results provided by different theoretical groups. The discussion of these theoretical uncertainties was early addressed by Chaboyer (1995) in a pioneering paper investigating the reliability of theoretical predictions concerning H-burning structures presently evolving in galactic globular clusters (GCs) and, in turn, on the accuracy of current predictions about GC ages. More recently, such an investigation has been extended to later phases of stellar evolution by Cassisi et al. (1998, hereinafter CCDW; 1999), and Castellani \\& Degl'Innocenti (1999), who discussed theoretical predictions concerning central He-burning low-mass stars populating the Horizontal Branch of galactic globular clusters. In this paper we will discuss predictions concerning the evolutionary behaviour of Asymptotic Giant Branch stars, devoting particular attention to two key observational parameters, such as the luminosity of the predicted AGB clump and the number ratio between HB and AGB stars N$_{\\mathrm AGB}$/N$_{\\mathrm HB}$. These parameters appear of particular relevance since the AGB clump luminosity has been proposed as an alternative distance indicator for old--intermediate age stellar populations (Pulone 1992), while the ratio N$_{\\mathrm AGB}$/N$_{\\mathrm HB}$ is an excellent tool for investigating the efficiency of mixing processes during the HB phase (Buonanno et al. 1985), being HB lifetimes extremely sensitive to the extension of the semiconvective region in the stellar core.\\\\ In the two next sections we will first discuss the differences between ``old\" and ``new\" models, testing the most updated theoretical scenario on selected high-quality Color-Magnitude (CM) diagrams of galactic GCs. Section 4 will deal with an investigation on the uncertainties still existing in current theoretical models. Concluding remarks will close the paper. ", + "conclusions": "In this paper we found that the ``new\" theoretical scenario arising from stellar models with updated physical inputs, appears able to account for the main observational constraints on GC AGB stars. However, we insist on the evidence that no theoretical result can be taken at its face value, because of still existing theoretical uncertainties. In this context we critically discussed the indetermination due either to the efficiency of some macroscopic mechanisms, like atomic diffusion and breathing pulses, or to the uncertainties on the input physics (equation of state, opacity, nuclear cross sections etc.). While the efficiency of atomic diffusion has little influence on He-burning models, the occurrence of breathing pulses is ruled out by the comparison between theory and observation, confirming the conclusion reached by Caputo et al. (1989) in the frame of the ``old\" theoretical scenario. Regarding the physical inputs, the main uncertainty still present in He-burning models is given by the $^{12}C(\\alpha,\\gamma)^{16}O$ cross section, which influences the He burning lifetimes and the C/O ratio in the carbon-oxygen core, with relevant consequences on the final cooling of white dwarfs. Since this nuclear reaction rate has a strong effect on the predicted central He burning lifetimes, it affects also the evaluation of the initial He abundance in galactic GCs via the $R$ parameter, i.e., the ratio between the HB stars and the RGB ones brighter than the HB (see e.g. the discussion in CCDW and in Zoccali et al. 2000). On the contrary, data in Table~2 show that AGB lifetimes appear marginally affected by this uncertainty; therefore one could be tempted to use the ratio N$_{\\mathrm AGB}$/N$_{\\mathrm HB}$ as an indicator of efficiency of the $^{12}$C +$\\alpha$ reaction. However, as discussed in the previous section, the extension of the mixed core of He-burning stars affects the value of N$_{\\mathrm AGB}$/N$_{\\mathrm HB}$ too, so that a comparison with a given CMD can only tell us if the combination of mixing prescription plus the adopted $^{12}$C +$\\alpha$ reaction rate are consistent with observations. We have already shown that, at least for M5, observational data are in agreement with the combination of the new reaction rate plus semiconvection without breathing pulses. However, CCDW have already shown that current theoretical predictions for the R parameter, as obtained by using the same input physics adopted in the present work, provide an unrealistically large value for the original He abundance in galactic GCs. This unpleasant situation could be clarified only by reducing the uncertainties related to the adopted physical inputs. To summarize, the cross section for $^{12}$C +$\\alpha$ reaction, the amount of central mixing in He burning stars, the evaluation of the original He in globular cluster stars and the ratio N$_{\\mathrm AGB}$/N$_{\\mathrm HB}$ are strongly connected. In this context, better evaluations of one or more out of the quoted quantities would be of great relevance to assess the problem on a more firm basis." + }, + "0112/astro-ph0112057_arXiv.txt": { + "abstract": "The ambient interstellar environment of wind- and supernova-driven superbubbles strongly affects their evolution, but its properties are rarely well-determined. We have therefore obtained \\hi\\ aperture synthesis imaging of the environment around three similar, optically-selected superbubble nebulae in the Large Magellanic Cloud. The resulting \\hi\\ maps show that the ambient gas distribution around these superbubbles differ to an extreme: DEM L25 shows no neutral shell component, but is nestled within an \\hi\\ hole; DEM L50 shows a massive neutral shell component, but is otherwise within an \\hi\\ void; and DEM L301 shows no correspondence at all between the optical nebula and \\hi\\ distribution. There is also poor correspondence between the \\hi\\ and optical kinematics. These results strongly caution against inferring properties of the ambient neutral environment of individual superbubbles without direct observations. Finally, all three objects show some evidence of shock activity. ", + "introduction": "The evolution of superbubbles created by the supernovae (SNe) and stellar winds of OB associations is a critical determinant of the structure and energetics of the interstellar medium (ISM). The coronal component of the ISM, or hot ionized medium (HIM) is believed to originate within such superbubbles and supernova remnants (SNRs; e.g., Cox \\& Smith 1974), hence the evolution of these structures determines whether and how this hot gas is released into the diffuse HIM. Models for the ISM are strongly dependent on whether the HIM is a pervasive, dominant component, as envisioned by e.g., McKee \\& Ostriker (1977), or whether it plays a lesser role as favored by, e.g., Slavin \\& Cox (1993), returning to a paradigm closer to the two-phase ISM of Field, Goldsmith, \\& Habing (1969). Superbubble evolution also directly impacts the relative volumes of enriched, coronal gas in galactic disks vs. halos (e.g., Shapiro \\& Field 1976). Similarly, superbubbles and their related filaments are thought to provide much of the structure for the diffuse ISM, including warm and cold \\hi\\ and the warm ionized medium (e.g., Kennicutt {\\etal}1995). Superbubble activity is therefore fundamental to the ISM phase balance, structure, and kinematics; and it is also a principal driver of galactic evolutionary processes. The standard model for the evolution of these superbubbles is an adiabatic, pressure-driven shell with continuous mechanical energy injection (e.g., Pikel'ner 1968; Castor, McCray, \\& Weaver 1975). The mechanical power $L$ is provided by stellar winds and SNe of the parent OB association. In this model, an outer shock sweeps up ambient ISM into a thin cooled shell, while an inner, reverse shock thermalizes the mechanical energy, thereby generating the hot interior whose pressure drives the shell growth. It is generally assumed that successive SNe can be approximated as a continuous energy injection, and that all the available SN energy is thermalized (e.g., Mac Low \\& McCray 1988). The standard, adiabatic model then yields simple analytic relations governing the shell evolution: \\begin{equation} R\\propto \\bigl(L/n\\bigr)^{1/5}\\ t^{3/5} \\quad , \\end{equation} \\begin{equation} v = \\Dt{R}\\propto \\bigl(L/n\\bigr)^{1/5}\\ t^{-2/5} \\quad , \\end{equation} where $R$ and $v$ are the shell radius and expansion velocity, respectively, $n$ is the ambient density, and $t$ is elapsed time. There are several lines of evidence supporting this standard model for superbubble evolution (see Oey 1999): a few young, stellar wind-dominated objects have well-constrained $L$ from classification of the stellar populations, and show reasonable consistency with the predicted kinematics (Oey 1996b). In addition, multiwavelength observations of highly active star-forming regions like 30 Doradus (Wang 1999) and the starburst galaxy NGC 5253 (Strickland \\& Stevens 1999) show hot gas morphologically contained within cooler shell structures. Higher ionization species like \\ion{C}{4} and \\ion{Si}{4} have been detected in lines of sight through all superbubbles with available observations (Chu {\\etal}1995), suggesting interface zones between hot and cold gas. In addition, the superbubble size distribution derived from the standard model is in remarkable agreement with observed size distributions of \\hi\\ shells in the Magellanic Clouds (Oey \\& Clarke 1997; Kim {\\etal}1999). However, other evidence suggests that superbubble evolution is often not nearly as simple as described by the standard model. The first problem is that the model overestimates the shell growth rate for young, wind-dominated superbubble nebulae (Oey 1996b; Garc\\'\\i a-Segura \\& Mac Low 1995; Drissen {\\etal}1995). This could be due to a systematic understimate in the ambient densities $n$ by up to a factor of $\\sim10$, or to a similar overestimate in the input power $L$ predicted by the assumed stellar mass-loss rates. Secondly, Oey (1996b) identified a class of superbubbles that exhibit expansion velocities that are several times higher than predicted, relative to their radii. Most of these also show anomalously bright X-ray emission, which suggests shell acceleration by SNR impacts (Chu \\& Mac Low 1990; Oey 1996b). In such cases, not all the available SN energy will be thermalized to power the shell expansion, thus the effect of discrete SNe on shell evolution remains unclear. Clarifying the structure of the ambient medium is clearly essential in determining the origin of these discrepancies. Since most estimates of the ambient $n$ for young nebular shells are based on the \\Ha\\ electron densities, the presence of a massive \\hi\\ shell component could imply a significant underestimate in assumed $n$. Another possibility is that the structure of the ambient medium is highly irregular, providing openings to a network of low-density channels. In that case, the superbubble could be ruptured in one or more locations, depressurizing the interior and releasing hot gas into these channels. Soft X-ray imaging with {\\it Einstein} and {\\it ROSAT} suggests this process in at least one object, DEM L152 (Chu {\\etal}1993; Magnier {\\etal}1995) in the Large Magellanic Cloud (LMC). These minor blowouts would then cause accelerated, high-velocity kinematics in excess of the standard model predictions (e.g., Oey \\& Smedley 1998; Mac Low {\\etal}1998). Expansion into the density gradient perpendicular to galactic disks offers an additional systematic explanation for anomalously fast shell expansion in face-on galaxies (Silich \\& Franco 1999). ", + "conclusions": "We have obtained \\hi\\ observations of the environment around three optically-selected superbubbles in the LMC. The optical nebulae have similar properties: all are roughly 50 pc in radius and show optical kinematic structures with anomalously high expansion velocities that contrast with values expected from the standard evolutionary model for wind and SN-generated superbubbles. We find that the \\hi\\ environment within roughly 500 pc of the objects vary to an extreme. DEM L25 shows no \\hi\\ associated with the superbubble itself, but appears to be nestled within an \\hi\\ hole, and the shell is interacting with an adjacent \\hi\\ cloud; DEM L50 shows an \\hi\\ component to the shell that is $\\sim 10$ times more massive than the ionized shell, but is otherwise in a large \\hi\\ void; and DEM L301 shows no correlation of any kind between the optical and \\hi\\ distributions. Furthermore, the high-velocity kinematics observed optically for all three objects are not seen in any of our \\hi\\ data. These findings strongly caution against inferring properties of the neutral interstellar environment around any individual superbubbles without direct observations. All three superbubbles have reported signatures suggestive of shock activity, including supersonic nebular velocities. DEM L25 and DEM L50 show elevated \\Ha\\ emission in excess of that predicted from photoionization by their classified stellar populations, and DEM L50 encompasses the known SNR N186~D on its northern edge. DEM L25 also shows elevated \\oiii/\\Ha\\ and all objects show high \\sii/\\Ha. The emission-line spectrum of DEM L301 was modeled in detail by Oey {\\etal}(2001), who found evidence of contributions by shock excitation. We examined our 1380 MHz (20 cm) continuum in conjunction with 843 MHz continuum observations from the MOST survey to obtain rough, spatially-resolved estimates of the spectral indices. These suggest some presence of a non-thermal component for at least DEM L25 and the SNR N186~D. We mentioned above that the new \\hi\\ data do not show the complex, high-velocity shell components that are evident in the optical nebular data for these objects. However, the uneven \\hi\\ distributions for DEM L25 and DEM L301 are consistent with the high ionization state of these nebulae, suggesting that these shells are fully ionized. Their patchy \\hi\\ remains consistent with the optical filaments corresponding to incipient blowout features or material accelerated by SNR impacts, as has been previously suggested. The lack of correspondence between the nebular and \\hi\\ kinematics is more problematic for DEM L50, although the dominant \\Ha\\ and \\hi\\ components are broadly consistent with each other. The \\hi\\ velocity profile suggests that the neutral component may correspond to the approaching side of the shell." + }, + "0112/astro-ph0112261_arXiv.txt": { + "abstract": "We discuss the X-ray emission observed from Supernova Remnant 1987A with the {\\it Chandra X-ray Observatory}. We analyze a high resolution spectrum obtained in 1999 October with the high energy transmission grating (HETG). From this spectrum we measure the strengths and an average profile of the observed X-ray lines. We also analyze a high signal-to-noise ratio CCD spectrum obtained in 2000 December. The good statistics ($\\approx 9250$ counts) of this spectrum and the high spatial resolution provided by the telescope allow us to perform spectroscopic analyses of different regions of the remnant. We discuss the relevant shock physics that can explain the observed X-ray emission. The X-ray spectra are well fit by plane parallel shock models with post-shock electron temperatures of $\\approx 2.6 \\keV$ and ionization ages of $\\approx 6 \\times 10^{10}$ cm$^{-3}$ s. The combined X-ray line profile has a FWHM of $\\approx 5000 \\kms$, indicating a blast wave speed of $\\approx 3500 \\kms$. At this speed, plasma with a mean post-shock temperature of $\\approx 17 \\keV$ is produced. This is direct evidence for incomplete electron-ion temperature equilibration behind the shock. Assuming this shock temperature, we constrain the amount of collisionless electron heating at the shock front at $T_{e0} / T_s = 0.11^{+0.02}_{-0.01}$. We find that the plasma has low metallicity (abundances are $\\approx 0.1$ -- $0.4$ solar) and is nitrogen enriched (N/O $\\approx 0.8$ by number), similar to abundances found for the equatorial ring. Analysis of the spectra from different regions of the remnant reveals slight differences in the parameters of the emitting plasma. The plasma is cooler near the optical Spot 1 (at position angle $\\approx 30^{\\circ}$) and in the eastern half of the remnant, where the bright optical spots are found, than in the western half, consistent with the presence of slower ($\\approx 500 \\kms$) shocks entering denser ring material. There is an overall flux asymmetry between the two halves, with the eastern half being 15 -- 50\\% brighter (depending on how the center of the remnant is defined). However, our spectroscopic analysis shows that $<5\\%$ of the overall X-ray emission could come from a slow shock component. Therefore the flux asymmetry cannot fully be due to X-rays produced by the blast wave entering the ring, but rather indicates an asymmetry in the global interaction with the circumstellar material interior to the ring. ", + "introduction": "Supernova (SN) 1987A provided the astrophysical community with a wealth of information about core collapse supernova (see e.g., \\citealt{arne89,mccr93}). Now, as the object enters its remnant phase, it will reveal still more about the physics of the explosion and of the progenitor star's history, as well as give insight into the processes occurring in other supernova remnants (SNR). The radiation from SNR 1987A is now dominated by the interaction of the SN debris with the circumstellar material (CSM) which was present before the star exploded. This interaction will convert the kinetic energy stored in the debris expansion into radiation which is observable from the radio to X-ray bands. The circumstellar gas is most obvious in the triple ring system seen in optical wavelengths \\citep{burr95} as it recombines following the initial ionizing flash from the SN \\citep{lund91}. The mechanism responsible for forming these rings is still uncertain. In one model the system is a result of the self interaction of the progenitor star's stellar winds \\citep{luo91,wang92,blon93}. The appearance of both radio \\citep{stav92,stav93} and X-ray \\citep{beue94,gore94,hasi96} emission from the remnant at day $\\approx 1000$ suggests that the supernova blast wave encountered an \\HII\\ region interior to the equatorial ring \\citep{chev95,bork97a}. This \\HII\\ region was probably formed by the evaporation of gas off of the ring by the blue supergiant progenitor star \\citep{chev95}. The hydrodynamics of the interaction is quite complex \\citep{chev92,bork97b} and depends on the density distributions of the expanding debris and the CSM. In any case, a double shock structure forms which consists of a blast wave which travels outwards into the CSM and a reverse shock which travels back (in a Lagrangian sense) into the SN ejecta (see e.g., \\citealt{chev82}). This structure is shown schematically in Figure~\\ref{fig-DSS}. Between these two shocks is a region of shocked circumstellar gas and a region of shocked ejecta gas, separated by a contact surface which is hydrodynamically unstable \\citep{chev92}. Emission from the reverse shock has been observed as high velocity ($\\sim \\pm 12,000 \\kms$) \\La\\ and \\Ha\\ emission \\citep{mich98} created as neutral hydrogen crosses the shock front. By modeling the line profiles \\citet{mich98} found that emission from the reverse shock front is confined primarily to the equatorial plane ($\\pm 30^{\\circ}$) and is located at $\\approx 75\\%$ of the distance from the SN to the inner surface of the equatorial ring. Images of SNR 1987A in radio bands \\citep{gaen97,manc01} show a toroidal geometry with an east-west brightness asymmetry. These images support the notion that the emission from the remnant is brightest in the equatorial regions. \\citet{manc01} also find that the radio image is expanding with a radial velocity of $\\approx 3500 \\kms$. This radio emission is most likely synchrotron emission from nonthermal particles which are accelerated in the shocks. While it is uncertain where this emission originates, it is reasonable to assume that it comes from the approximate region between the blast wave and reverse shock. If so, the expansion speed of the radio remnant is a good approximation for the speed of the blast wave in the \\HII\\ region. Hydrodynamic models for the interaction of the SN 1987A's blast wave with an equatorial \\HII\\ region produce blast wave velocities of $\\approx 4100 \\kms$ \\citep{bork97a}, close to that inferred from the radio images. Such a shock can produce X-ray emitting plasma with electron temperatures of several keV. X-ray images taken by the {\\it Chandra X-ray Observatory} ({\\it Chandra}) were shell-like and confirmed that the emission was indeed from the shocks rather than emission from a central source \\citep{burr00,park01}. The X-ray spectrum is characterized by electron temperatures of $\\approx 3 \\keV$. \\citet{park01} showed that the hard band X-ray image ($E = 1.2$ -- $8.0 \\keV$) of the remnant looks very similar to the radio image, which suggests that the double shock structure is the source of emission in both bands. X-ray images from {\\it Chandra} also suggested an expansion rate of $5200 \\pm 2100 \\kms$ \\citep{park01}, consistent with the radio image expansion rate and the theoretical models. One of the young remnant's most exciting developments in the last several years has been the appearance in the optical and ultraviolet (UV) of several rapidly brightening unresolved spots on the equatorial ring \\citep{lawr00,mich00}. These spots appear where the blast wave has encountered dense ($n > 10^4 \\pcc$) obstacles on the inside of the optical ring. In these obstacles the blast wave has slowed considerably (to $\\sim 250 \\kms$) and cooled sufficiently to emit copious UV and optical radiation \\citep{pun02}. These radiative shocks are too slow to produce substantial X-ray emission. There may, however, be lower density obstacles near the optical spots into which shocks with speeds of $\\approx 500$ -- $1000 \\kms$ are transmitted. Such shocks would be copious producers of soft X-ray emission ($E \\lesssim 1 \\keV$) but they would not produce significant amounts of optical or UV line emission. The soft band X-ray image ($E = 0.3$ -- $0.8 \\keV$) presented by \\citet{park01} shows X-ray brightening at approximately the same locations as the optical spots, suggesting that intermediate velocity shocks are present in their vicinity. Here we present and analyze in more detail two X-ray data sets obtained with {\\it Chandra}: a high spectral resolution dispersed spectrum (obtained with the HETG grating) first shown in \\citet{burr00}; and spatially resolved CCD spectra obtained from the high signal-to-noise ratio (S/N) imaging observation presented by \\citet{park01}. We present the observations and reductions in \\S\\ref{sec-obs}. In \\S\\ref{sec-physics} we present our shock models to describe the observed X-ray emission. In \\S\\ref{sec-analCCD} we analyze the CCD spectrum to determine the overall shock parameters and their variations across the remnant. We also consider spectroscopic evidence for the existence of slow X-ray emitting shocks. In \\S\\ref{sec-analdisp} we present the X-ray line fluxes obtained from the dispersed spectrum. These data confirm the results of the CCD spectral analysis and the legitimacy of the shock model for the emission. In \\S\\ref{sec-lineprof} we present the combined X-ray line profile and compare it to simulated profiles in order to establish the speed of the blast wave. In \\S\\ref{sec-disc} we discuss our findings and what they tell us about the remnant's structure and development, as well as what we may learn from future investigations. We conclude in \\S\\ref{sec-summary} with a summary of our results. ", + "conclusions": "\\label{sec-disc} The blast wave velocity implied by the width of the X-ray line profile indicates that the shock temperature should be $\\approx 17 \\keV$. This is much higher than the electron temperature measured in our fits to the X-ray spectra ($T_e \\approx 2.5 \\keV$). Though not surprising given the discussion in \\S\\ref{sec-physics}, this result is direct observational evidence of the existence of unequal electron and ion temperatures behind a fast shock ($\\beta = T_{e0} / T_s < 1$). Up to now we have used models in which the post-shock electron temperature is constant. In order to confirm the validity of this assumption, and to constrain the amount of electron heating at the shock front itself, we make use of the {\\it npshock} models which account for the equilibration of electron and ion temperatures through Coulomb collisions. We fit the CCD spectrum of the entire remnant with the abundances and interstellar absorption given in \\S\\ref{sec-analCCD}. Since line excitation is caused primarily by collisions with electrons, the X-ray spectrum is insensitive to the ion temperatures, and thus to the shock speed. Therefore the parameter $T_s$ (and thus $\\beta$) is poorly constrained by our spectral analysis. With this in mind, we perform various fits each having a fixed value of $\\beta$, fitting for the initial electron temperature $T_{e0}$ and ionization age. Fits with $\\beta = 0.05, 0.1$, and $0.2$ give results of comparable statistical significance ($\\chi^2_\\nu$ of 1.13, 1.05, and 1.07, with $\\nu = 69$) with initial electron temperatures of $1.6, 1.7$, and $2.0 \\keV$ respectively, and ionization ages of $\\approx 5 \\times 10^{10}$ cm$^{-3}$ s. Note that all of the fitted initial electron temperatures are lower than the value ($2.5 \\keV$) found in the constant electron temperature fits. This is true because the temperature determined in the {\\it pshock} model fit represents the mean electron temperature of the plasma. An {\\it npshock} model fit with a lower value for $\\beta$ has a higher shock temperature, thus a faster heating rate, and so requires a lower initial electron temperature in order to produce a plasma with the same mean electron temperature. That all of the $\\beta$ model fits have similar $\\chi^2_\\nu$ values indicates that the spectrum is most sensitive to the mean electron temperature and less so to the actual temperature distribution. This is primarily due to the fact that the electron temperature remains nearly constant over the volume behind the shock (it varies by at most a factor of $\\approx 2$ in the $\\beta = 0.05$ model). This result supports our use of the constant electron temperature models as representative models in our analysis in \\S\\ref{sec-anal}. The CCD spectrum alone cannot constrain the value of $\\beta$ since it is insensitive to the shock temperature. However, we note that the $\\beta = 0.1$ model seems the most realistic since its shock temperature, $T_s = 17.3 \\keV$, is very close to what is expected from a shock with velocity of $\\approx 3500 \\kms$. Indeed we can use the shock velocity results of \\S\\ref{sec-lineprof} and fix the shock temperature at $16.5 \\keV$ in order to constrain $\\beta$. For this fit we find a good fit, $\\chi^2_\\nu = 1.04$ $(\\nu = 70)$ with $T_{e0} = 1.8^{+0.4}_{-0.1} \\keV$ or $\\beta = 0.11^{+0.02}_{-0.01}$. It is worth mentioning that models where $T_{e0}$ is fixed at a very low value (as would be the case if no collisionless electron heating occurred at the shock front) are not consistent with the data. This amount of equilibration for a $\\approx 3500 \\kms$ shock is consistent with the results of \\citet{ghav01}, who found that $\\beta$ decreases with increasing shock strength. However, the observed electron temperature is not low enough to imply that particle acceleration has taken a significant amount of energy from the plasma, as found in the blast wave of 1E 0102.2-7219 \\citep{hugh00}. The X-ray image (as well as images of the radio source) show a noticeable brightness asymmetry between the eastern and western halves of the remnant. Our analysis in \\S\\ref{sec-analCCD} quantified the X-ray brightness asymmetry as the eastern half having 15\\% more flux (in the 0.3 -- 6 keV band) than the western half. However, this result is based on our choice of the image centroid as the center of the X-ray image. Clearly the level of flux asymmetry depends on this choice. To explore the possible extent of the asymmetry, we performed a similar spectral analysis with the center defined from an ellipse which was fit to the shape of the X-ray image (see Figure~\\ref{fig-regions}). In this case the asymmetry between the X-ray fluxes of the eastern and the western halves increases to $\\approx 50\\%$. We fit one- and two-shock models to the spectra in a similar manner as described in \\S\\ref{sec-analCCD} and find similar physical parameters at about the same statistical significance. Unfortunately, absolute X-ray astrometry was not possible with our data so we are unable to precisely constrain the level of asymmetry any better than 15 -- 50\\%. The presence of slower X-ray emitting shocks where the blast wave has struck the inner ring (representing $<5\\%$ of the total X-ray flux) cannot account for the observed level of flux asymmetry. Therefore, this asymmetry must be due primarily to a global asymmetry in the interaction of the SN blast wave with the circumstellar environment. The blast wave must be interacting with denser material in the east. The luminosity of the shocked plasma is approximately proportional to $n_0^2T_s^{-0.6}V$, where $n_0$ is the pre-shock density, $V (\\propto v_s)$ is the volume of shocked gas, and $T_s (\\propto v_s^2)$ is the shock temperature. Since $v_s \\propto n_0^{-1/2}$, the luminosity goes as $n_0^{2.1}$. Whatever level of brightness asymmetry that exists is a good estimate of the differences in pre-shock density that the blast wave is encountering. The blast wave could be encountering denser material in the east either because the CSM is asymmetric or because the SN explosion is asymmetric. If the SN explosion (and therefore the expanding debris) was symmetric then the CSM must have greater density toward the east then toward the west. In this case the blast wave would propagate more slowly toward the east than toward the west. If, on the other hand, the CSM is symmetric, the blast wave must have propagated further toward the east to encounter the denser gas closer to the equatorial ring. For this to occur, the driving pressure of the SN ejecta must be greater in the east than in the west, indicating an asymmetric SN explosion. It seems reasonable that the SN occurred at the center of the ring, in which case the appearance of the optical spots primarily in the east favors the latter hypothesis. Unfortunately, we cannot confirm this hypothesis with the present X-ray data, which lacks sufficiently accurate astrometric registration with respect to the SN or the optical ring. This is not a trivial task due to the scarcity of X-ray point sources in the vicinity of SN 1987A. However, accurate X-ray astrometry would yield powerful scientific rewards. Short of this, the location of the reverse shock is a good proxy for the location of the blast wave and may provide the best constraints on the cause of the observed asymmetry (Michael et al. 2002, in preparation). Many young SNR have been observed to contain nonthermal X-ray emission (e.g., \\citealt{alle99}). From an analysis of {\\it XMM-Newton} data \\citet{asch01} have suggested that the spectrum of SNR 1987A includes such a component. We explore this possibility with the {\\it Chandra} CCD spectrum by considering a composite model consisting of a single shock and a power-law component. Our best fit, with $\\chi^2_\\nu = 1.06$ $(\\nu = 60)$, is not statistically superior (based on the F-test) to our previous models. The photon index of the power-law emission is $\\Gamma = 2.2\\pm0.3$. The shock has $T_e = 1.4 \\pm 0.2 \\keV$, $n_et \\approx 10^{11}$ cm$^{-3}$ s, and has abundances which are greater by factors of $\\approx 2$ -- 3 compared to the values derived previously. The observed X-ray flux/luminosity of the power-law component is about 57\\% of the total in the 0.3 -- 6 keV range and $\\sim 75\\%$ in the 1.4 -- 8 keV range. The upper limit to a possible point source of X-ray emission in SNR 1987A is $\\approx 13\\%$ in the 2 -- 10 keV band \\citep{park01}. Therefore, an embedded pulsar or its wind nebula cannot be responsible for such a nonthermal component. On the other hand, the fast shocks themselves might produce X-ray synchrotron radiation if they could accelerate electrons to sufficiently high energies. In the case of the most efficient diffusive shock acceleration of particles (e.g., \\citealt{joki87}), when the shock front is parallel to the magnetic field, we estimate that a magnetic field of strength $\\gtrsim 3 \\times 10^{-4}$ G must be present in order for the shocks to accelerate electrons to sufficiently high energies within the age of the remnant. With the success of the shock models at fitting the {\\it Chandra} spectra, we find that a power-law component is not required by the data. Moreover, the abundances required for such a fit seem unreasonably high. However, the present spectra do not contain significant flux above a few keV and therefore the presence of a hard power-law tail can not be ruled out. The possible presence of a narrow component to the X-ray line profile is intriguing although at present this result is not very statistically significant. Confirmation of a narrow component would support the existence of X-ray emission from lower velocity shocks entering the ring. As the remnant gets brighter and the blast wave penetrates more of the ring, the X-rays produced by the slower shocks will contribute an increasing percentage of the total X-ray luminosity \\citep{bork97b}. One manifestation of this will be a growing narrow component in the line profile. Moreover, since the efficiency of a shock at producing a given line depends on its velocity, different lines are likely to have significantly different profiles. For example the \\SiKa\\ and \\SiLa\\ lines are formed only in the faster shocks while the \\OLa\\ line will have a significant contribution from slower shocks. Thus we would expect the Si lines to appear broader. Detection of differences in the line profiles will help resolve the different shock components present in the remnant. The overall spectrum will also change as slower shocks begin to dominate, adding a softer emission component. The structure of SNR 1987A and its emission is changing significantly on a time scale of less than a year, therefore much can be learned about the development of the remnant and its shocks through continued observations of the remnant in the X-ray band. In conclusion, we summarize our findings: \\begin{itemize} \\item{The X-ray spectrum (including the resolved line fluxes) is well represented by an adiabatic, plane parallel shock model with mean electron temperature of $\\approx 2.6 \\keV$ and ionization age of $\\approx 6 \\times 10^{10}$ cm$^{-3}$ s. We find elemental abundances that are subsolar ($\\approx 0.2$), with an enhanced nitrogen abundance (N/O $\\approx 0.8$ by number).} \\item{There are indications that a second shock component (with $T_e \\approx 0.4 \\keV$ and $n_e t \\sim 10^{13}$ cm$^{-3}$ s) is present in the eastern half of the remnant. This slower shock component (with $v_s \\approx 500 \\kms$) may be related to the shocks ($v_s \\lesssim 250 \\kms$) that produce the optical spots on the equatorial ring. Currently, the slow shock can account for only $\\approx 4\\%$ of the total X-ray emission from the remnant.} \\item{By investigating simulated profiles based on simple models for the blast wave's geometry, we infer from the observed X-ray line profile that the blast wave has a velocity of $3400 \\pm 700 \\kms$. Future observations of X-ray line profiles should show different widths for different lines and increasing flux in a narrow component.} \\item{The measured electron temperature in the spectra and the width of the line profile provide direct evidence for incomplete electron-ion temperature equilibration behind the blast wave. Assuming the shock temperature indicated by the line profile, we measure the level of collisionless electron heating at the shock front at $T_{e0} / T_s = 0.11^{+0.02}_{-0.01}$.} \\item{We constrain the level of east-west brightness asymmetry at 15 -- 50\\%. This asymmetry is caused by the blast wave interacting with denser material in the east. A better constraint will only be possible if a more precise astrometric registration of the location of the SN is undertaken for the X-ray image. Such a registration will also help determine if the source of the asymmetry is an asymmetric circumstellar environment or an asymmetric SN explosion.} \\item{Continued observations in the X-ray band of this rapidly brightening and changing remnant will help us better understand SN 1987A, its circumstellar environment, and their interaction.} \\end{itemize}" + }, + "0112/astro-ph0112507_arXiv.txt": { + "abstract": "{We suggest a two-dimensional wavelet devised to deduce the large-scale structure of a physical field (e.g., the Galactic magnetic field) from its integrals along straight paths from irregularly spaced data points to a fixed interior point (the observer). The method can be applied to the analysis of pulsar rotation and dispersion measures in terms of the large-scale Galactic magnetic field and electron density. The method does not use any \\emph{\\'a priori\\/} assumptions about the physical field and can be considered as an algorithm of wavelet differentiation. We argue that a certain combination of the wavelet transformation with model fitting would be most efficient in the interpretation of the available pulsar $\\RM$ data. ", + "introduction": "Tomography, understood broadly, is a reconstruction of a multidimensional physical field from its integral projections obtained by exposing it in different aspects. Tomography problems often occur in astronomy, especially radio astronomy, where the intervening medium is transparent and the observable quantities represent certain integrals along the line of sight. An important example is the studies of the magneto-ionic medium in the Milky Way using the Faraday rotation measure $\\RM$ and dispersion measure $\\DM$ of pulsars -- both are integrals along the line of sight involving interstellar magnetic field and thermal electron density. A peculiar feature of the astronomical tomography is that just one vantage point is available, as the observer is confined to the close vicinity of the Sun. Astronomical tomography then focuses on extended objects which can be meaningfully analyzed using the integral projections obtained in a variety of directions. In this paper we suggest a method to assess the spatial structure of the global Galactic magnetic field using pulsar $\\RM$ (and also $\\DM$ as the Faraday rotation measure depends on the thermal electron density), one of the most informative tracers of the large-scale component of the magnetic field \\citep[see e.g.,][]{Beck2000}. A fundamental problem here is that the derivation of magnetic field $\\vec{\\HH}$ from the observable integral \\begin{equation} \\RM \\equiv K\\int_0^r \\nel {\\vec{\\HH}}\\cdot d{\\vec{s}}\\;, \\label{RM} \\end{equation} in fact involves differentiation of the observed quantity with respect to $r$ resulting in a catastrophic amplification of noise. Moreover, the distances to pulsars, $r$, are contaminates by large errors, both random and systematic. The problem is further aggravated by the fact that the pulsars (and extragalactic radio sources), that provide the $\\RM$ data, are irregularly spaced and cover the sky very inhomogeneously. Here $\\nel$ is the number density of thermal electrons, $r$ is the distance to the radio source and $K=0.81\\,{\\rm rad\\,m^{-2}\\,cm^3\\,\\mu G^{-1}\\,pc^{-1}}$; the positive direction of $\\vec{s}$ is that towards the observer, so that a field directed towards the observer produces positive $\\RM$. The thermal electron density can be obtained from the dispersion measure of pulsars \\citep{manch72,manch74}: \\begin{equation} \\DM\\equiv\\int_0^r \\nel \\,d s\\,, \\label{DM} \\end{equation} and this involves similar complications. The idea of this paper is twofold. Firstly, we use wavelets to filter out the noise in the $\\RM$ data resulting from small-scale fluctuations in the magneto-ionic medium and from the irregular spacing of the data points. Secondly, we introduce a new, specialized family of wavelets devised to avoid (or at least to minimize) noise amplification in dealing with an observable represented by an integral along the line of sight, given that the available data probe a range of distances into the field localization region. The Faraday rotation and dispersion measures of pulsars are suitable observables for such an analysis. The plan of the paper is as follows. In Sect.~\\ref{data_s} we discuss the available samples of pulsar $\\RM$ and $\\DM$ and methods used to obtain $\\vec{B}$ and $\\nel$ from them. Section~\\ref{wl_s} introduces the basic ideas of our method, which can be applied in various contexts. The algorithm is described in Sect.~\\ref{method_s}, and the advantages and limitations of the method are discussed in Sect.~\\ref{demo_s}. The efficiency of the method when applied to the pulsar $\\RM$ is discussed in Sect.~\\ref{diss_s}. ", + "conclusions": "\\label{diss_s} We have introduced a new wavelet devised for the analysis of the pulsar Faraday rotation measures in terms of the large-scale magnetic field (or of any other observable that is an integral of the quantity studied, e.g., the pulsar dispersion measures). This is a tomography approach because the field is directly reconstructed from its integral estimator. The method works well with data given on a regular mesh or scattered randomly but with gaps between the data points not exceeding $a/2$, with $a$ the scale of the wavelet. The separation of pulsars with known $\\RM$ exceeds this limit beyond about 3\\,kpc from the Sun. An advantage of the method is that it involves neither {\\it ad hoc\\/} assumptions about magnetic field structure nor model fitting. However, the wavelet transforms on the data grid of the pulsar catalogues available appear rather confusing and difficult to interpret. The advantages of the wavelet analysis and model fitting can be combined in a single approach applied by \\cite{mw} to the Faraday rotation measures of extragalactic sources. Instead of fitting a magnetic field model to noisy $\\RM$ data, these authors fitted the wavelet transform of the model $\\RM$ to the wavelet transform of the observed $\\RM$ with smaller scales (mainly responsible for the nose) filtered out. This has resulted in a significant improvement in the quality of the analysis. The application of the wavelet introduced here will allow us to fit the wavelet transforms of the magnetic field derived from $\\RM$ to the model. We expect that will result in a significant improvement of the results against the usual procedure of fitting model $\\RM$ to the observed noisy data. Another way to improve the results would be to analyze simultaneously the Faraday rotation measures of pulsars and extragalactic radio sources. Since many extragalactic sources occur at high Galactic latitudes, a plausible assumption about the vertical distribution of the magneto-ionic medium will have to be adopted." + }, + "0112/astro-ph0112441_arXiv.txt": { + "abstract": "{Weak gravitational lensing by the large scale structure can be used to probe the dark matter distribution in the Universe directly and thus to probe cosmological models. The recent detection of cosmic shear by several groups has demonstrated the feasibility of this new mode of observational cosmology. In the currently most extensive analysis of cosmic shear, it was found that the shear field contains unexpected modes, so-called B-modes, which are thought to be unaccountable for by lensing. B-modes can in principle be generated by an intrinsic alignment of galaxies from which the shear is measured, or may signify some remaining systematics in the data reduction and analysis. In this paper we show that B-modes in fact {\\it are produced} by lensing itself. The effect comes about through the clustering of source galaxies, which in particular implies an angular separation-dependent clustering in redshift. After presenting the theory of the decomposition of a general shear field into E- and B-modes, we calculate their respective power spectra and correlation functions for a clustered source distribution. Numerical and analytical estimates of the relative strength of these two modes show that the resulting B-mode is very small on angular scales larger than a few arcminutes, but its relative contribution rises quickly towards smaller angular scales, with comparable power in both modes at a few arcseconds. The relevance of this effect with regard to the current cosmic shear surveys is discussed; it can not account for the apparent detection of a B-mode contribution on large angular scales in the cosmic shear analysis of van Waerbeke et al.\\ (2002). ", + "introduction": "Gravitational lensing by the large-scale structure (LSS) leads to the distortion of the images of distant galaxies, owing to the tidal gravitational field of the matter inhomogeneities. Following very early work on the study of light propagation in an inhomogeneous universe (e.g., Gunn 1967; Kantowski 1969), Blandford et al. (1991), Miralda-Escude (1991) and Kaiser (1992) have pointed out that the observation of this `cosmic shear' effect immediately yields information about the statistical properties of the LSS and, thus, on cosmology. Non-linear evolution of the matter spectrum was taken into account in later analytical (e.g., Jain \\& Seljak 1997; Bernardeau et al. 1997, Kaiser 1998; Schneider et al. 1998, hereafter SvWJK) and numerical (e.g., van Waerbeke et al. 1999; Jain et al. 2000, White \\& Hu 2000) studies; see Mellier (1999) and Bartelmann \\& Schneider (2001; hereafter BS01) for recent reviews. It was only in 2000 when four teams nearly simultaneously and independently announced the first detections of cosmic shear from wide-field imaging data (Bacon et al. 2000; Kaiser et al. 2000; van Waerbeke et al. 2000; Wittman et al. 2000). The detections reported in these papers (and in Maoli et al. 2001, using the VLT, and Rhodes et al. 2001, using HST images obtained with the WFPC2 camera) concerned various two-point statistics, like the shear dispersion in an aperture, or the shear correlation function. In van Waerbeke et al. (2001), the aforementioned statistics, as well as the aperture mass statistics (SvWJK), were inferred from the effective 6.5 square degrees of high-quality imaging data. Very recently, H\\\"ammerle et al.\\ (2002) reported on a cosmic shear detection using HST parallel images taken with the STIS instrument on an effective angular scale of $\\sim 30''$. The shear field, originating from the inhomogeneous matter distribution, is a two-dimensional quantity, whereas the projected density field of the matter is a scalar field. The relation between the shear $\\gamma(\\vc\\theta)=\\gamma_1(\\vc\\theta)+{\\rm i}\\gamma_2(\\vc\\theta)$ and the projected matter density $\\kappa(\\vc\\theta)$ is \\be \\gamma(\\vc\\theta)={1\\over\\pi}\\int_{\\Real^2}\\d^2\\theta'\\; \\D(\\vc\\theta-\\vc\\theta')\\kappa(\\vc\\theta')\\;, \\label{eq:1.1} \\ee with the kernel \\be \\D(\\vc\\theta)={\\theta_2^2-\\theta_1^2-2{\\rm i}\\theta_1\\theta_2 \\over \\abs{\\vc\\theta}^4}\\;; \\label{eq:1.2} \\ee here, $\\kappa$ is the dimensionless surface mass density, i.e., the physical surface mass density divided by the `critical' surface mass density, as usual in gravitational lensing; we follow the notation of BS01 in this paper. Since the two shear components originate from a single scalar field, they are related to each other; in particular, their partial derivatives should satisfy compatability relations, as we shall discuss in Sect.\\ts 2 below. In analogy with the polarization of the CMB, a shear field satisfying these compatability relations is called an E-mode shear field. Pen et al. (2002) pointed out that the cosmic shear data of van Waerbeke et al. (2001) contains not only an E-mode, but also a statistically significant B-mode contribution in addition. Such B-modes can be generated by effects unrelated to gravitational lensing, such as intrinsic alignment of galaxies (e.g., Heavens et al.\\ 2000, Crittenden et al.\\ 2001a; Croft \\& Metzler 2000; Catelan et al.\\ 2000) or remaining systematics in the data reduction and analysis. In this paper we show that a B-mode contribution to the cosmic shear is obtained by lensing itself. A B-mode is generated owing to the clustering properties of the faint galaxies from which the shear is measured. This spatial clustering implies an angular separation-dependent clustering in redshift, which is the origin not only of the B-mode of the shear, but also of an additional E-mode contribution. The paper is organized as follows: in Sect.\\ts 2 we provide a tutorial description of the E/B-mode decomposition of a shear field. Most of the results there were derived before in Crittenden et al.\\ (2001b, hereafter C01), but we formulate them in standard lensing notation, which will be needed for the later investigation. The calculation of two-point cosmic shear statistics in the presence of source clustering is presented in Sect.\\ts 3 where it is shown that this clustering produces a B-mode. Numerical and analytical estimates of the amplitude of this B-mode are provided in Sect.\\ts 4 and discussed in Sect.\\ts 5. ", + "conclusions": "We have shown that the clustering of galaxies from which the shear is measured leads to the presence of a B-mode in the cosmic shear field, in addition to providing an additional component to the E-mode. The reason for this effect in essence is the angular separation-dependent redshift correlation of galaxies, which causes the mean of the product of the angular-diameter distance ratio along two lines-of-sight not to factorize, but to depend on $\\theta$. For a fiducial model considered in detail, the B-mode contribution amounts to more than $\\sim 2\\%$ on angular scales below $1'$ (or $\\ell \\ga 2.16 \\times 10^4$), and its relative importance quickly rises towards smaller angles. On substantially larger angular scales, however, the B-mode contribution is small. Furthermore, the additional E-mode contribution is very similar in size to the B-mode power, which will allow an approximate correction of the measured E-mode for this additional term. From an observational point-of-view, the most easily accessible quantities are the shear correlation functions $\\xi_\\pm$, as one can easily deal with gaps in the data field. In Sect.\\ts 2 we have given explicit relations regarding how other two-point statistics of the shear can be calculated in terms of the shear correlation function. The finite support of the functions $T_\\pm$ indicates that the aperture measures are more easily obtained from observational data than either the E- and B-mode correlation functions, or the E- and B-mode shear dispersions. Therefore, the aperture measures are the preferred method to check for the presence of a B-mode contribution in the shear data. We have varied some of the model parameters; in particular, we have considered the case of lower mean source redshift (corresponding to a brighter flux threshold), and simultaneously increasing $A(1')$, such that the clustering length $r_0$ stays about the same. In this case we found a very similar ratio between the B- and E-mode power spectra as for the example considered in Sect.\\ts 4. We consider it unlikely that the observed B-mode in the present day data sets is due to the source clustering effect. The B-mode found in van Waerbeke et al.\\ts (2001) and Pen et al.\\ (2002) can actually be used to search quantitatively for residual systematics. Its detection in van Waerbeke et al.\\ts (2001) was done by obtaining $M_{\\rm ap}$ and $M_\\perp$ by laying down a grid of circular apertures on the data field. A more accurate measurement of $\\ave{M_{\\rm ap}^2}$ and $\\ave{M_{\\perp}^2}$ has been obtained from the same data by Pen et al.\\ (2002), by calculating them from the observed correlation functions $\\xi_\\pm$, as in (\\ref{eq:N29}). In fact a subsequent analysis revealed that the B-mode measured in these data were essentially residual systematics caused by an overcorrection of the PSF, and can be corrected for (van Waerbeke et al.\\ 2002). In this latter analysis, no significant B-modes are detected at small angular scales, but on scales above $\\sim 10'$, slightly significant values of $M_\\perp$ are detected; the effect discussed in this paper can certainly not account for them. The effect considered here seems to have been overlooked hitherto. Bernardeau (1998) considered the effects of source clustering on cosmic shear statistics and concluded that this source clustering can strongly affect the skewness and kurtosis of the cosmic shear, but to first order leaves the shear dispersion (and thus the power spectrum) unaffected. Hamana et al.\\ (2002) studied this effect with ray tracing simulations, again concentrating on the skewness. Most of the other ray tracing simulations of weak lensing (e.g., van Waerbeke et al.\\ 1999; Jain et al.\\ 2000) assumed all sources to be at the same redshift, in which case the additional power discussed here does not occur. Lombardi et al.\\ (2002) calculated the effect of source clustering on the noise of weak lensing mass maps, showing that it can provide a significant noise contribution in the inner regions of clusters. It must be pointed out that the effect considered here is unrelated to other lensing effects which in principle could generate a B-mode, such as lens-lens coupling or the break-down of the Born approximation (see Bernardeau et al.\\ 1997 and SvWJK for a discussion of these two effects on the skewness). Numerical estimates (e.g., Jain et al. 2000) show that these latter two effects are very weak. Bertin \\& Lombardi (2001) considered the situation of lensing by two mass concentrations along the line-of-sight, where a B-mode is generated by a strong lens-lens coupling, but the fraction of lines-of-sight where this occurs is tiny. Another effect which could in principle generate a B-mode from lensing is the fact that the observable is not the shear itself, but the reduced shear (Schneider \\& Seitz 1995). In the appendix we show that this effect is indeed negligible. Like the intrinsic alignment of galaxies, which can yield a spurious contribution to the measured cosmic shear, the source clustering effect can in principle be avoided if redshift estimates of the source galaxies are available. In that case, by estimating the shear correlation function, pairs of galaxies with a large likelihood to be at the same distance can be neglected. In contrast to the intrinsic correlation of galaxies, the B-mode from source clustering appears to be fairly insensitive to the redshift distribution of the source galaxies, provided the clustering length is kept fixed. \\appendix" + }, + "0112/astro-ph0112488_arXiv.txt": { + "abstract": "Some sources of systematic errors, specific to polarized CMB measurements using bolometers, are examined. Although the evaluations we show have been made in the context of the Planck mission (and more specifically the Planck HFI), many of our conclusions are valid for other experiments as well. ", + "introduction": "CMB fluctuations are difficult to measure because of their extreme weakness. Systematics must therefore be very well controlled. This is even more true for polarization anisotropies, which are expected to be less than 10\\% of the temperature fluctuations.\\\\ Low frequency noise is a source of troubles for polarization as well as for temperature measurements. However, as the polarized signal depends on the direction of the detector projected in the sky, its suppression, or ``destriping'' requires a specific treatment. The way in which polarized destriping can be implemented for Planck is outlined in the next section.\\\\ A few definitions useful to the following discussion are given in the third section.\\\\ Polarizers are never perfect: the unwanted polarization is never totally suppressed and the polarizer direction is never perfectly known. The impact of these uncertainties is evaluated in section four.\\\\ In the fifth section, we discuss the crucial difficulties linked to signal differences: calibration, pointing and beam mismatches between different detectors.\\\\ We then consider the question of avoiding elliptical error boxes on the Stokes parameters, which might be confused with a polarization signal when the signal to noise ratio is small. \\\\ A few concluding remarks are given in the last section. ", + "conclusions": "In an experiment such as Planck HFI, where polarization measurements are made with detectors sensitive to the total polarized intensity in one direction, the main source of systematic error in polarization measurements is the fact that $Q$ and $U$ Stokes parameters are obtained from signal differences. This can be overcome and even turned into an advantage if the systematics are common to both detectors with the same size and therefore disappear in the difference. OMT's and PSB's are a partial answer to this requirement as the two detectors have nearly the same lobes and pointings. However the electronic chains (and part of the optics for OMT's) are different. Moreover one still has to combine the signals of two different feeds to get the full polarized information. This latter combination is less dangerous however, as it does not involve intensity differences, and therefore will not generate polarisation where there is none. A rotating polarizing device in front of {\\em one} bolometer in {\\em one} feed and read by {\\em one} electronic chain may provide a solution to these difficulties, but one has to check that it does not bring new systematics and it may be difficult to implement on satellite or balloon borne experiments." + }, + "0112/astro-ph0112394_arXiv.txt": { + "abstract": "I review the current status of cosmology as emerging from recent observations of cosmic microwave background anisotropies as well as from other sources of cosmological information. ", + "introduction": "The widely accepted paradigm for cosmology is the hot Big Bang model. In this framework, the geometry and evolution of the Universe is defined by its matter and energy content through general relativity theory. The Universe is expanding, so that it was hotter and denser at earlier times. The rate of expansion is quantified by the Hubble parameter $H$, whose present value $H_0$ is parameterized by the quantity $h$ as $H_0=100\\;h$ km~s$^{-1}$~Mpc$^{-1}$. The amount of matter and energy in the Universe from different components (baryons, dark matter, radiation, vacuum energy, etc.) is parameterized by the quantities $\\Omega_{(i)}\\equiv\\rho_{(i)}/\\rho_c$. The critical density, $\\rho_c=1.88\\times 10^{-29}\\;h^2$~g~cm$^{-3}$, is defined in such a way that $\\Omega\\equiv\\sum_i \\Omega_{(i)}=1$ for a Universe with flat geometry (while $\\Omega<1$ and $\\Omega>1$ for open and closed geometry respectively). An additional ingredient of the standard cosmological model is {\\em inflation}\\cite{inflation}, a phase of early superluminal expansion of the Universe required to solve some problems of the Big Bang model. Inflation makes some well-defined predictions. First of all, the geometry of the Universe has to be very close to flat. Second, the structure we observe today in the Universe was produced by gravitational amplification of primordial density perturbations generated during inflation, characterized by having a nearly scale-invariant spectrum and by being Gaussian distributed. While until recent times the knowledge of the parameters of the cosmological model was plagued by large uncertainties, the situation has now dramatically changed. Cosmology is not a data-starved science anymore. In the past few years, high-quality observations have fueled an impressive progress in our understanding of the Universe. We have entered the epoch of high precision cosmology. Recent results from observation of the CMB temperature anisotropy have allowed us to constrain most cosmological parameters to unprecedented accuracy, giving for the first time a robust determination of the total energy density (and in turn of the geometry) of the Universe. In addition, a whole set of new observations of the large-scale structure properties of the Universe have put the determination of the mean matter density in the Universe on a firm ground. Finally, measurements of distant Type Ia Supernovae have recently provided evidence that the Universe has just entered a phase of accelerated expansion. In the following I will review the emerging scenario, giving particular emphasis to CMB as a cosmological probe. ", + "conclusions": "" + }, + "0112/hep-ph0112160_arXiv.txt": { + "abstract": "We point out that solar neutrino oscillations with large mixing angle as evidenced in current solar neutrino data have a strong impact on strategies for diagnosing collapse-driven supernova (SN) through neutrino observations. Such oscillations induce a significant deformation of the energy spectra of neutrinos, thereby allowing us to obtain otherwise inaccessible features of SN neutrino spectra. We demonstrate that one can determine temperatures and luminosities of non-electron flavor neutrinos by observing $\\bar{\\nu}_{e}$ from galactic SN in massive water Cherenkov detectors by the charged current reactions on protons. ", + "introduction": "The historical observation of neutrinos from the supernova (SN) 1987A at the Kamiokande \\cite{Kam} and IMB \\cite{IMB} detectors had a great impact and confirmed the basic picture of stellar collapse and SN explosion~\\cite{janka}. However, the number of observed events was too small to draw definite conclusions about the explosion mechanism or detailed properties of SN neutrinos. Hopefully, the next SN in our galaxy will be detected by several massive neutrino detectors, currently operating or planned. Such observations of neutrinos from a galactic SN should provide us a far more detailed information on properties of SN neutrinos, as well as the newly formed neutron star, thus bringing substantial progress to our understanding of stellar collapse. The discovery of atmospheric neutrino oscillations at Super-Kamiokande (SK)~\\cite{SKatm} implies that neutrinos have masses and different neutrino flavors mix, and it is supported by the results from the first long-baseline neutrino oscillation experiment, K2K \\cite{K2K}. Moreover, a clear evidence for solar neutrino oscillations into mu/tau neutrinos has been obtained by combining the Sudbury Neutrino Observatory (SNO) charged current (CC) measurement \\cite{SNO1} with elastic scattering measurement at SK \\cite{SK} and the most recent {\\it in situ} CC and neutral current (NC) measurements at SNO \\cite{SNO2}. Neutrino oscillations add a new ``complication'' to the diagnostics of SN neutrinos, since it is no longer true that a neutrino $\\nu_\\alpha$ leaving the neutrinosphere with definite flavor $\\alpha$ will be detected on Earth as the same neutrino species. Then, the task of extracting information such as the original of neutrino spectra in the SN core from terrestrial observations, requires solving an ``inverse problem''. In this letter, we point out that rather than a ``complication'' the neutrino oscillation provides a new powerful tool for probing otherwise inaccessible features of SN neutrino spectra. This holds if the mixing angle responsible for the solar neutrino oscillation is large \\cite{solaranalysis} as clearly indicated by the current solar neutrino data. We show that a high statistics observation of $\\bar{\\nu}_e$ through the CC reaction $\\bar{\\nu}_e + p \\to n + e^+$ will enable us to extract not only the original temperature of $\\bar{\\nu}_e$ but also that of $\\bar{\\nu}_{\\mu(\\tau)}$ at the neutrinosphere, as well as their time integrated luminosities. In order to determine these parameters we employ a $\\chi^2$ method to ``separate'' two different neutrino spectra at the neutrinosphere with different temperatures and luminosities. Hereafter we use a collective notation $\\nu_{x}$ ($\\bar{\\nu}_{x}$) for $\\nu_{\\mu(\\tau)}$ ($\\bar{\\nu}_{\\mu(\\tau)}$) because they cannot be distinguished by their physical properties inside the SN. We focus on $\\bar{\\nu}_e$ observation for the following reasons: (i) this is expected to be the channel with highest statistics in a large water Cerenkov detectors such as SK, Hyper-Kamiokande (HK) or UNO detectors, which are under consideration~\\cite{onemegaton}. Neutrinos from a galactic SN at 10 kpc would produce 7000-10000 events at SK and 2-3 $\\times 10^5$ events at HK~\\cite{totsuka}. (ii) the large solar neutrino mixing angle necessarily implies that the $\\bar{\\nu}_e$ spectrum observed at the Earth is a strong mixture of two originally different SN neutrino spectra of $\\bar{\\nu}_e$ and $\\bar{\\nu}_{x}$. This fact has been used to derive various constraints on neutrino mixing parameters ~\\cite{SSB94,MN00,Kachelriess:2001sg}. Complementary information may be obtained also through direct detection of $\\bar{\\nu}_{x}$ through the NC reaction either at SNO or KamLAND~\\cite{SNO/KLAND}. In addition to generic features of SN neutrino spectra obtained in SN simulations~\\cite{janka,MayleWilson,SNsimu,JH89,Raffelt} we will take into account some new features indicated by recent studies. Most importantly, they include a new parameter which characterize the departure from the equipartition of integrated luminosities to electron and other flavor neutrinos, the quantity of greatest uncertainty in SN simulations. (See below.) ", + "conclusions": "We show in Fig. \\ref{FvsNF1} the $3~\\sigma$ C.L. allowed parameter region obtained in the $\\langle E_{\\bar{\\nu}_e} \\rangle - \\tau_E$ and $\\langle E_{\\bar{\\nu}_e} \\rangle - E_b$ planes for SK and HK detectors by assuming the LMA solution, taking the best-fit value for mixing angle, $\\tan^2{\\theta} =$ 0.42~\\cite{afterSNO}. We also present results for the no-oscillation case for comparison. We set the initial astrophysical parameters as $\\langle E_{\\bar{\\nu}_e}\\rangle^0 = 15$ MeV, $\\tau_E^0$ = 1.4, $E_b^0 = 3 \\times 10^{53}$ erg, and $\\xi^0$ = 0.5. Fig. \\ref{FvsNF1} demonstrates that we can extract the both $\\bar{\\nu}_e$ and $\\bar{\\nu}_x$ temperatures in the presence of large mixing oscillations. The accuracies we can achieve with the LMA best-fit parameters are $\\Delta \\tau_E/\\tau_E \\sim $ 9 (1.5) \\%, $\\Delta \\langle E_{\\bar{\\nu}_e} \\rangle/ \\langle E_{\\bar{\\nu}_e} \\rangle \\sim $ 4 (1) In contrast, in the absence of oscillation there is no sensitivity in $E_b$ due to the inability of determining $\\nu_x$ and $\\bar{\\nu}_x$ luminosities. The improvement in the accuracy of the determination of these quantities for the oscillation case is remarkable, especially with HK. Let us now examine how accurately we can determine the equipartition-violation parameter $\\xi$ and its correlation with $\\tau_E$ and $E_b$. This is shown in Fig. \\ref{FvsNF2}. We note that the accuracy of $\\tau_E$ determination is remarkably good in spite of the rather poor accuracy in the knowledge of $\\xi$. It should be noticed that Fig.~\\ref{FvsNF2} demonstrates that HK can do a much better job for $\\xi$, $\\Delta \\xi/\\xi \\sim $ 10 \\%. We emphasize that without having sensitivity to $\\xi$ we can not determine $E_b$ in a good accuracy, as they are strongly correlated (See the right panel of Fig.~\\ref{FvsNF2}.) We finally note that HK's enormous sensitivity $\\sim 10$ \\% may mean one could potentially examine the problem by treating accreting and thermal phase separately under the present approximations. \\begin{figure}[h] \\vspace*{2cm} \\special{psfile=Eavtau.eps angle=0 hscale=45 vscale=55 hoffset=-10 voffset=-180} \\special{psfile=EavEb.eps angle=0 hscale=45 vscale=55 hoffset=220 voffset=-180} \\vspace{6.0cm} \\caption[]{ Extracting the astrophysical parameters. The figure shows 3 $\\sigma$ contours assuming $\\langle E_{\\bar{\\nu}_e} \\rangle^0 =15$ MeV, $\\tau_E^0$ = 1.4, $E_b^0 = 3 \\times 10^{53}$ erg, and $\\xi^0$ = 0.5 for the SK and HK detectors assuming the LMA solution to the solar neutrino problem. Best fits are indicated by the stars. } \\label{FvsNF1} \\vglue 0.5cm \\end{figure} The negative correlation between $\\xi$ and $\\tau_E$ can be understood due to the fact that the effect of lowering $\\xi$, which implies a decrease of the relative contribution of $\\bar{\\nu}_x$, can be compensated by increasing $\\tau_E$ to keep the higher energy tail of the observed spectrum similar. On the other hand, the strong positive correlation between $\\xi$ and $E_b$ just reflects the relationship $E_b = 2 (1 + 2 \\xi) E^{tot}_{\\bar{\\nu}_{e}}$, which is valid under our approximation $E^{tot}_{\\nu_{e}} = E^{tot}_{\\bar{\\nu}_{e}}$. The correlation is robust and exists with and without oscillations as indicated in Fig.~2. \\begin{figure}[h] \\vspace*{3.0cm} \\special{psfile=Xitau.eps angle=0 hscale=45 vscale=55 hoffset=-10 voffset=-180} \\special{psfile=XiEb.eps angle=0 hscale=45 vscale=55 hoffset=220 voffset=-180} \\vspace{6.0cm} \\caption[]{ Determination of the non-equipartition parameter $\\xi$ and its correlations with $\\tau_E$ and $E_b$. Same assumptions as in Fig. \\ref{FvsNF1}. } \\label{FvsNF2} \\vglue 0.5cm \\end{figure} We have also examined the case of a smaller value of the initial $\\bar{\\nu}_x$ energy, $\\tau_E^0$ = 1.25 to verify the robustness of our results. We found that although the sensitivity to $E_b$ determination worsens by a factor of $\\sim$ 2, we still have a reasonable sensitivity to $\\tau_E$, $\\Delta \\tau_E/\\tau_E \\sim $ 14 \\% (2.5 \\% for HK), excluding the case of $\\tau_E$ = 1 at 3 $\\sigma$. See Ref.~\\cite{MNTV} for details. We have verified that our results do not change very much even if we relax our assumptions by taking into account deviations from power-law density profile of progenitor star, or from the precise best-fit values of solar neutrino mixing parameters we have adopted. The inclusion of possible Earth matter effects for a given experiment will imply some regeneration effect. Although this will somewhat weaken our results, the effect is rather small in practice. Additional details related to the robustness of our method can be found in Ref.~\\cite{MNTV}. In summary, we have suggested a simple but powerful way of extracting separately the temperatures of both $\\bar{\\nu}_e$ and $\\bar{\\nu}_{\\mu(\\tau)}$ as well as their integrated luminosities by analyzing $\\bar{\\nu}_e$ events that would be recorded by massive water Cherenkov detectors in the event of a galactic supernova explosion. In particular, an extraordinary power of megaton-class detectors (Hyper-Kamiokande or UNO) are noticed with regard to the determination of the integrated luminosities of $\\bar{\\nu}_e$ and $\\bar{\\nu}_{\\mu(\\tau)}$ as independent fit parameters. We stress that the large mixing between $\\nu_e$ and $\\nu_{\\mu(\\tau)}$, which is clearly indicated by the current solar neutrino data, is essential in determining temperature as well as luminosity of $\\bar{\\nu}_{\\mu(\\tau)}$, and this must have more profound implications which await further investigation. \\vskip 0.8cm \\noindent Note added: When the first version of this paper was ready for submission to the electronic archive we became aware of the paper by Barger {\\it et al.} \\cite {BMW01} who pursued the similar strategy as ours. However, they do not treat the violation of equipartition of energies as a fit parameter." + }, + "0112/astro-ph0112380_arXiv.txt": { + "abstract": "Nonspherical collapse is modelled, under the Zeldovich approximation, by six-dimensional random walks of the initial deformation tensor field. The collapse boundary adopted here is a slightly-modified version of that proposed by Chiueh and Lee (2001). Not only the mass function agrees with the fitting formula of Sheth and Tormen (1999), but the bias function and conditional mass function constructed by this model are also found to agree reasonably well with the simulation results of Jing (1998) and Somerville et al. (2000), respectively. In particular, by introducing a small mass gap, we find a fitting formula for the conditional mass function, which works well even at small time intervals between parent and progenitor halos during the merging history. ", + "introduction": "In a hierarchical model, the gravitationally bound objects of the universe formed from mergers of smaller objects into larger ones. An issue of fundamental importance is to understand how gravitational structures, such as galaxies and clusters, are distributed as a function of mass and evolving time. Press \\& Schechter (1974) provided a simple analytical model (hereafter PS model) to describe the mass distribution of dark halos (mass function). This theory is based on the assumption that the process of gravitational collapse can be approximated by spherical symmetry and it occurs when the mean density contrast of precollapse regions predicted by the linear theory reaches a threshold value $\\delta_{c}$. Bond et al. (1991) introduced the excursion set formalism and extended the PS theory to calculations of the conditional mass function, which is closely related to the issue of the linear bias and halo merging histories. Mo \\& White (1996) used the extended Press-Schechter theory (hereafter EPS) to address the large-scale bias, which concerns how the collapsed halos trace the underlying dark matter. Moreover, the EPS theory is also useful for construction of the merger history of dark halos (Bower 1991; Lacey \\& Cole 1993, 1994 hereafter LC93, LC94), which concerns the probability that a halo of mass $M_{2}$ at time $t_{2}$ has a progenitor halo of mass $M_{1}$ at earlier time $t_{1}$. The PS or EPS theory has been tested against N-body simulations (Efstathiou et al. 1988; LC94), and the comparisons suggest that the PS model overestimates the number of halos in low-mass regimes and underestimates it in high-mass regimes. This discrepancy indicates that the spherical collapse model used by the PS theory may be too simplistic. Many researchers have taken the nonspherical collapse process into account to correct this problem (e.g., Bond \\& Myers 1996; Monaco 1995; Lee \\& Shandarin 1998; Chiueh \\& Lee 2001; Sheth, Mo and Tormen 2001). Most of these works are still based on the main idea of the PS theory, and further incorporate either the dynamics of ellipsoids or the Zeldovich approximation (Zeldovich 1970). The main aims of this paper are to investigate the large-scale bias and conditional mass function by using the nonspherical collapse boundary (hereafter NCB) proposed by Chiueh and Lee (2001). We also test whether the results derived from nonspherical collapse boundary can provide better statistical predictions of dark halos than the EPS model does. The outline of the paper is as follows. In $\\S$ 2, we summerize the PS and EPS theory, and briefly review the Zeldovich approximation and the nonspherical collapse boundary proposed by Chiueh \\& Lee (2001). In $\\S$ 3, we show the algorithm and the results of mass function and large-scale bias of our model. In $\\S$ 4, the conditional mass function is derived and the comparison with N-body simulation is presented. In $\\S$ 5, the conclusion is given. ", + "conclusions": "In this paper, we have attempted to show that the nonspherical collapse boundary (NCB) reproduces the mass function, bias factor and conditional mass function better than the Press-Schechter theory. This model successfully addresses these subjects over a wide range of mass scales, while the (extend) Press-Schechter theory always fails in some mass range. Through the simulations of 6-dimensional random walks, we have shown that the bias function generated by NCB agrees with the N-body simulations reasonably well. Because NCB is more complicated than the PS's collapse boundary, the relation between the bias function and the conditional mass function is no longer as simple as the EPS model. Especially, the two issues are not of identical representation in the framework of NCB model. We provide a formula for the conditional mass function of the nonspherical collapse boundary, which can be applied to any two redshifts. In order to overcome the problem that random-walks fail to describe the merger history at small look-back time, we add a minor correction in the definition of variable $\\mu'$ (c.f. Eq. (\\ref{newmu})) while maintaining the overall mathematical form intact. This formula can describe the conditional mass probability more accurately than the EPS formalism over a wide mass range, even when the two epochs are close. All the results we obtained from NCB suggest that NCB is sufficiently reasonable to reproduce the statistical properties of dark halos. There are two interesting issues pertinent to NCB: one is its natural association with the halo angular momentum through its dependence on the quantity $r$. The halo angular momentum distribution predicted by the NCB model is reported elsewhere (Chiueh, Lee and Lin, in preparation). Second, as mentioned in $\\S$ 4, the conditional mass function offers an useful tool to construct the merger history through merger trees. Recent works have indicated that the merger history constructed by using EPS's conditional mass function deviates from that extracted from N-body simulations, and it is thought to be caused mainly by the spherical-symmetry assumption of the EPS model rather than the specific features of the merger tree scheme. Thus, we expect the agreement can be greatly improved by the NCB conditional mass function given in this work." + }, + "0112/astro-ph0112349_arXiv.txt": { + "abstract": "Taking advantage of recent HST data for field stars in the region of the Galactic globular cluster NGC~6397, we tested the predictions of several Galactic models with star counts reaching a largely unexplored range of magnitudes, down to $V$$\\sim$ 26.5. After updating the input stellar ($V-I$) colors, we found that the two-component Bahcall-Soneira (B\\&S) model can be put in satisfactory agreement with observations for suitable choices of disk/spheroid luminosity functions. However if one assumes the Gould, Bahcall and Flynn (1996, 1997) disk luminosity function (LF) together with the Gould, Flynn and Bahcall (1998) spheroid LF, there is no way to reconcile the predicted and observed $V\\,$-magnitude distribution. We also analysed the agreement between observed and predicted magnitude and color distributions for two selected models with a thick disk component. Even in this case there are suitable combinations of model parameters and faint magnitude LFs which can give a reasonable agreement with observational star counts both in magnitude and in color, the above quoted combination of Gould et al. (1997, 1998) LFs giving again predictions in clear disagreement with observations. ", + "introduction": "Since the very beginning of modern astronomy, the distribution of stars over the night sky has been understood as evidence for the spatial distribution we now know as ``the Galaxy\". Therefore, the use of star counts to constrain the Galactic structure has been the goal of generations of astronomers. In the long run, relevant progresses on that matter has been allowed by the increased amount of observational data, by the availability of modern computers and, last but not least, by the original approach to the problem presented twenty years ago in a seminal paper by Bahcall \\& Soneira (1980, see also 1984). Since that time star counts began an effective way of investigating the broad properties of stellar populations in our Galaxy (see e.g. Gilmore 1981, 1984, Gilmore \\& Reid 1983 and for a review Bahcall 1986 and Gilmore, Wyse \\& Kuijken 1989). Galactic models, which predicts star counts in selected areas of the sky and for selected intervals of magnitude, give, at the same time, constraints about such relevant issues as the amount of mass in the form of stars and, in particular, about the contribution of faint stars to the dark-matter problem. One naturally expects that the contribution of intrinsically faint stars increases as the limiting magnitude of star counts is increased. Thus the limiting magnitude of the observational samples is a critic factor governing the possibility of constraining the Galactic abundance of such kind of stars. However, the difficulty in distinguishing faint images of stars from galaxies has limited most of the ground based work to V $\\lcu$ 20, reaching as extreme limit V$\\approx$ 22 (e.g. Bahcall 1986, Reid \\& Majewski 1993, Haywood et al. 1997 and references therein). This situation has been recently improved by observational data from the Hubble Space Telescope (HST). HST observations reach very faint magnitudes (V$\\approx$30, see e.g. Williams et al. 1996) and even though the number of stars in the field of view is generally quite small they can be used to test Galactic models at the best resolution available. However the tests of Galactic models at faint magnitudes available in the literature are still very few and mainly based on a restricted number of stars (see Mendez et al. 1996, Reid et al. 1996, Basilio et al. 1996, Mendez \\& Guzman 1998). In this paper we will take advantage of recent HST observations of the Galactic globular cluster NGC~6397, to discuss a rich sample of about thousand bona fide Galactic field stars which should be fairly complete (more than 95\\% detection level) down to $V \\sim$ 25, with a reliable evaluation of the completeness reaching $V \\sim$ 27 (King et al. 1998, hereafter KACP). This sample obviously gives the exciting opportunity of testing Galactic models over an almost unexplored range of magnitudes, with possible relevant outcomes concerning the abundance of faint stars. For the analysis we will use observational data of field stars in the direction of NGC 6397 from the second epoch HST observations. Field stars have been obtained by an accurate separation from cluster stars based on proper motions analysis (KACP). The area covered by observations is of 6.6 arcmin$^2$ centered at the Galactic coordinates $l=337^\\circ.9$, $b=-11^\\circ.7$. Figure \\ref{CMD} shows the $(V-I, V)$ color--magnitude diagram for the about 1000 stars whose proper motion is significantly different from that of the NGC~6397 stellar population (cf.\\ Fig.\\ref{CMD} in KACP). In Section 2 we will use these data to discuss the predictions of the Bahcall-Soneira Galaxy model (Bahcall 1986) which has been widely used by observers to predict the number of stars of different types, colors and apparent magnitude ranges in different fields of interest (see e.g. Boeshaar \\& Tyson 1985, Flynn \\& Freeman 1993, Lasker et al. 1990, Stuwe 1990, Basilio et al. 1996). Bahcall and Soneira assumed that the Galaxy is simply composed by a disk, with a scale height of a few hundred parsecs, and a more or less spherical spheroid. The Bahcall-Soneira model, as further improved mainly to account for interstellar reddening and obscuration, was proved to be in good agreement with available observations (see e.g Bahcall 1986, Gould, Bahcall \\& Maoz 1993 and references therein) and the export code was made available on the Web to the scientific community in 1995. However at faint magnitudes such an agreement is strongly dependent on the assumption about the luminosity function (LF) of faint main sequence (MS) stars. As a result, in Sect.3, we will show that the above quoted agreement vanishes for several recent evaluations of the actual LF of either the disk or the halo components. The need for more complex Galactic models has been advanced in 1983 by Gilmore \\& Reid, who firstly suggested the occurrence of an extended ``thick disk'' population consisting of stars with spatial and kinematic properties intermediate between disk and spheroid. Starcounts is not the best way to constrain this third component because similar results could be obtained either with a thick disk or without it in dependence of the adopted model parameters (disk/thick disk scale height, scale lenght and normalizations). However detailed analysis which, in addition to the star counts, take into account the velocity distribution of local stars, support the need for a thick disk component (Ojha et al. 1996, Wyse \\& Gilmore 1995, Norris \\& Ryan 1991, Casertano, Ratnatunga \\& Bahcall 1990.) However up to now, sensitively different values for the thick disk structural parameters (density law, local density etc..) have been proposed (see e.g. Reid \\& Majewski 1993, Yamagata \\& Yoshii 1992, Ojha et al. 1996, Mendez \\& Guzman 1998). Thus the situation is not yet well defined and in Sect.4 we will use our data to test some three-component models currently available in the literature. In this context, one has to remind that with very faint observations it should be possible to resolve the Galaxy into its component populations. This because for V$>$19, at least for high Galactic latitudes, the various components are expected to contribute with different colors: the bluer colors are dominated by spheroid stars, the reddest objects do belong to the disk population, while the intermediate colors should be dominated by thick disk stars (see Basilio et al. 1996). Unfortunately, our data refer to a sky region of somewhat low latitude, which gives a sensitive gain in the number of stars but which - as we will find - makes difficult a clear separation among the various populations. As a result we will find that either two-component or three component models give an equally satisfactory agreement with magnitude and color observational counts. No one of the current Galactic models appears in agreement with observations for some extreme assumptions on the faint LF, whereas the data do not allow to discriminate between different three-component models within the current uncertainties on the scale height and on the density of stars in the solar neighbourhood. \\begin{figure} \\label{CMD} \\centerline{\\epsfxsize= 7 cm \\epsfbox{MA840f1.eps}} \\caption{The $(V,V-I)$ CM diagram for the sample of field stars studied.} \\end{figure} ", + "conclusions": "In this paper we have discussed star counts in the field of the globular NGC 6397. As expected, we found that the predicted distribution of stars fainter than M$_V$$\\approx$24 is critically dependent on the assumptions about the LF of faint MS stars. According to the discussion given in the previous sections one finds that neither the two-component B\\&S model not the more sophisticated three-component models by HR\\&C and G\\&R can satisfactory fit star counts in the field of NGC 6397, when assuming for disk and spheroid the low LF suggested by Gould et al. (1997), Gould et al. (1998), respectively. A satisfactory fit of the data can be achieved, either for a two-component or for a three component model, with other suitable combinations of LFs suggested in the current literature, provided that density and scale height of the two discussed three-component models are decreased within the predicted existing uncertainties. Nor the color distributions appear able to discriminate among the various possible solutions. However, we regard the previous conclusions as suggestions to be further tested in deep surveys at different Galactic locations. In this context, and before closing the paper, one has to notice that from the comparison of Figs.\\ 2 and 6 one derives the tantalizing evidence that going deeper than $V$=26.5 by a couple of magnitudes one would derive much more stringent and precise constraints on the LF of faint stars and, in turn, on the mean metallicity of the spheroid. Such an evidence should be taken into account in future researches on this matter." + }, + "0112/astro-ph0112455_arXiv.txt": { + "abstract": "We present optical spectroscopic identifications of hard X-ray (5-10 keV) selected sources belonging to the HELLAS sample obtained with {\\it BeppoSAX} down to a 5-10 keV flux limit of $f_{5-10 keV}$$\\sim$$3\\times 10^{-14}$ \\ecs. The sample consists of 118 sources. 25 sources have been identified trough correlations with catalogues of known sources. 49 have been searched for spectroscopic identification at the telescope. 13 fields resulted empty down to R=21. 37 sources have been identified with type 1 AGN and 9 with type 2 AGN. The remaining are: 5 narrow emission line galaxies, 6 Clusters, 2 BL Lac, 1 Radio Galaxy and 1 Star. Combining these objects with other hard X-ray selected AGNs from {\\it ASCA} and {\\it HEAO1}, we find that the local luminosity function of type 1 AGN in the 2-10 keV band is fairly well represented by a double-power-law-function. There is evidence for significant cosmological evolution according to a pure luminosity evolution (PLE) model $L_X(z)$$\\propto$$(1+z)^k$, with $k$=2.12 and $k$=2.22 ($\\sigma_k$$\\simeq$$0.14$) in a ($\\Omega_m$,$\\Omega_\\lambda$)=(1.0,0.0) and in a ($\\Omega_m$,$\\Omega_\\lambda$)=(0.3,0.7) cosmology, respectively. The data show an excess of faint high redshift type 1 AGN which is well modeled by a luminosity dependent density evolution (LDDE), similarly to what observed in the soft X-rays. However, in both cosmologies, the statistic is not significant enough to distinguish between the PLE and LDDE models. The fitted models imply a contribution of AGN1 to the 2-10 keV X-ray background from 35\\% up to 60\\%. ", + "introduction": "\\label{secintro} AGN have first been discovered in the radio and soon after searched in the optical band. Consequently, they have been classified using their optical characteristics and mainly divided into two categories: type 1 (AGN1) and 2 (AGN2) according to the presence or not of broad emission lines in their optical spectra (we will keep this definition of AGN1 throughout this paper). Before the advent of the last generation of hard X-ray telescopes, AGN samples where predominantly based on AGN1 selected either in the optical or, later on, in the soft X-rays by {\\it Einstein} and {\\it ROSAT}. In these bands the evolution of AGN1 has been well measured (see e.g. Della Ceca et al. 1992; Boyle et al. 2000; Miyaji, Hasinger, \\& Schmidt 2000). On the contrary the production of samples of AGN2 has been difficult at any wavelength and limited to few local surveys. The general picture was in favor of a model in which AGN1 objects were associated to AGN with low absorption in the hard X-rays while AGN2 to obscured sources with large column densities and spectra strongly depressed in the soft X-rays, as expected in the unification models (e.g. Antonucci 1993). In the last decade the advent of the {\\it ASCA} and {\\it BeppoSAX} satellites has allowed for the first time the detection and identification of AGN as the main counterparts of hard (2-10 keV) X-ray sources down to fluxes $\\sim 5\\times 10^{-14}$ \\ecs, more than 2 orders of magnitude fainter than {\\it HEAO1} (Wood et al. 1984). These identifications accounted for about 30\\% of the 2-10 keV hard X-ray background (Ueda et al. 1998; Fiore et al. 1999). Recently the new generation of X-ray satellites such as {\\it Chandra} and {\\it XMM-Newton}, have reached fluxes 100 times fainter, identifying hundreds of sources and almost resolving the hard (2-10 keV) X-ray background (e.g. Mushotzky et al. 2000; Fiore et al. 2000; Giacconi et al. 2001; Hornschemeier et al. 2001; Hasinger et al. 2001; Tozzi et al. 2001; Baldi et al. 2001). Thanks to their excellent angular resolution ($\\sim$1-5$''$), the first spectroscopic identifications projects have been able to observe faint (I$\\sim$23) optical counterparts. At variance with the classical type-1/type-2 model in the optical, a significant number of the counterparts ($\\sim$30\\%) resulted to be apparently optical normal galaxies, with X-ray luminosities $L_X$$\\simeq$$10^{42}-10^{44}$ erg s$^{-1}$ typical of AGN activity, and moreover part of the optical type 1 AGNs resulted to be absorbed in the hard X-rays (see e.g. Fiore et al. 2000; Barger et al. 2001; Tozzi et al. 2001; Hornschemeier et al. 2001; Comastri et al. 2002). These observations have complicated the picture of the AGN model. In this framework the computation of the density of AGN has become an even more difficult task. In fact, it is not clear how to classify the sources and to take into account the selection biases introduced by the observation in the 2-10 keV range, where the absorption still play a relevant role. \\begin{inlinefigure} \\psfig{figure=f01.eps,angle=0,width=8.1cm} \\caption{ The flux distribution of the 118 sources from HELLAS used in this analysis (continuous line). The dashed line is the distribution of the 74 sources for which a spectroscopic identification campaign has been carried out ({\\it empty} fields included). The distribution of the 13 sources for which we end up with an {\\it empty} field is represented by the hatched histogram. } \\figurenum{1} \\addtolength{\\baselineskip}{5pt} \\end{inlinefigure} These recent deep surveys with {\\it Chandra} and {\\it XMM-Newton} have reached fluxes $\\sim$ $5\\times 10^{-16}$ \\ecs (2-10 keV) in quite small areas (less than 1 deg$^2$). As a consequence these surveys are not able to provide statistical significant samples at brighter fluxes ($\\sim 10^{-13}$ \\ecs; 5-10 keV) where the density of sources is about 5/deg$^2$ (Fiore et al. 2001) and tens of square degrees are to be covered. Such data are necessary to provide large numbers of spectroscopic identified sources in a wide range of X-ray fluxes in order to cover as much as possible the $L_X/z$ plane and hence to derive their X-ray luminosity function (LF). In this paper we report the results of the spectroscopic identifications of one of such brighter samples. The X-ray sources have been detected by the {\\it BeppoSAX}-MECS instruments in the 5-10 keV band in the framework of the High Energy LLarge Area Survey (HELLAS). Preliminary results have been presented in Fiore et al. (1999) and La Franca et al. (2001). The whole survey and the catalogue is described by Fiore et al. (2001). The data have been analyzed in the framework of the synthesis models for the X-ray background by Comastri et al. (2001), and the correlation with the soft X-rays has been investigated by Vignali et al. (2001). In section 2 we describe our X-ray and optical observations. In section 3 we present an analysis of the evolution of AGN in the 2-10 keV band. Because of the reasons previously described, the selection and definition of type 2/absorbed sources is still not clear, and thus we restricted our evolutionary studies to type 1 AGN only. The results are discussed in section 4. \\placefigure{f1} ", + "conclusions": "Thanks to the identification of 61 sources of the HELLAS sample we have been able to double the number of hard X-ray AGN1 available for statistical analysis at fluxes in the range $f_{2-10 keV}\\sim 10^{-13.5}$--$10^{-12}$ \\ecs. In total we can use 74 AGN1 at these fluxes (37 from HELLAS), which combined with the local sample of Grossan have allowed to show directly the shape of the LF of AGN1 as function of redshift and measure its evolution. The PLE models provide satisfactory fits of the data both in the ($\\Omega_m$,$\\Omega_\\lambda$)=(1.0,0.0) and in the ($\\Omega_m$,$\\Omega_\\lambda$)=(0.3,0.7) cosmologies. Our estimate of the LF in the ($\\Omega_m$,$\\Omega_\\lambda$)=(1.0,0.0) has a significantly larger normalization in comparison to the previous measure from Boyle et al. (1998). The data start to probe in the hard X-rays the faint part of the LF where the excess of density of AGN1 has been observed in the soft X-rays, justifying the implementation of the LDDE models. However, in both cosmologies, the statistic is not significant enough to distinguish between the PLE and LDDE models (see Table 6). In fact, in the prediction of the differential counts of AGN1 shown in Figure 11, the models differentiates at fluxes fainter than $f_{2-10keV}\\sim10^{-13}$ \\ecs, where the statistic is still poor. The new upcoming fainter surveys from {\\it Chandra} and {\\it XMM-Newton} will easily test which model is correct. \\begin{inlinefigure} \\figurenum{12} \\psfig{figure=f12.eps,angle=-90,width=8.1cm} \\caption{ The evolution of the density of type 1 AGN having LogL$_X$(2-10 keV)$>$44.8 as a function of redshift. The soft X-ray data from Miyaji et al. (2000) have been over-plotted assuming a slope $\\alpha=0.6$, which corresponds to a limit LogL$_X$(0.5-2 keV)$>$44.5. The continuous line are the predictions of our LDDE model. } \\addtolength{\\baselineskip}{5pt} \\end{inlinefigure} In Table 6 the percentages of the contribution to the 2-10 keV X-ray background are shown. The X-ray background has been computed integrating the LF up to $z=3.5$ for $L_X$$>$10$^{42}$ $erg/s$. At variance with the results of 24\\% obtained from the evolution of the LF derived by Boyle et al. (1998), our models reproduces from $\\sim$35\\% up to $\\sim$60\\% of the XRB. We used $I_{2-10}= 1.95\\times 10^{-11}$ erg cm$^{-2}$ s$^{-1}$ deg$^{-2}$ from Chen, Fabian and Gendrau (1997). The highest percentages would probably imply that part of the absorbed population necessary to reproduce the XRB is already included in the AGN1 at high redshift. However, more detailed analysis of this issue are beyond the scope of this paper. The percentages would decreas if a value f $I_{2-10}=2.35 \\times 10^{-11}$ erg cm$^{-2}$ s$^{-1}$ deg$^{-2}$ from Vecchi et al. (1999) is assumed. It is interesting to notice that AGN1 in the 2-10 keV range show an evolution up to $z\\sim2$ which is fairly well compatible with what observed in the soft X-rays. In fact we found a good fit of the data in the ($\\Omega_m$,$\\Omega_\\lambda$)=(1.0,0.0) Universe if we assume exactly the same parameters found in the soft X-rays by Miyaji et al. (2000) for AGN1, by only looking for a fit with the break luminosity $L_\\ast$ (model 7 in Table 6). Miyaji et al. (2000) found Log$L_\\ast$=43.78. The value $L_\\ast$ found by our fit in the 2-10 keV band is Log$L_\\ast$=44.13. This luminosity difference implies an X-ray spectrum for AGN1 with slope $\\alpha$=0.6, which is the same slope used in our computation of the LF. We note however that Miyaji et al. (2000) used a slope $\\alpha$=1.0 in their computations of the soft X-ray LF. Therefore, taken at a face value, the match between the soft and hard X-rays LFs implies that AGN1 have a broad band concave spectrum getting steeper going toward lower energies. As already discussed in the previous section, the ($\\Omega_m$,$\\Omega_\\lambda$)=(0.3,0.7) cosmology for AGN1 alone has not been analyzed by Miyaji et al. (2000). The agreement of the evolution measured in the soft and hard X-rays is also shown in Figure 12, where the evolution of the density of AGN1 brighter than LogL$_{2-10 keV}$=44.8 is shown. The soft X-ray data from Miyaji et al. (2000) have been over-plotted assuming a slope $\\alpha=0.6$, which corresponds to a limit LogL$_X$(0.5-2 keV)$>$44.5. The continuous line are the predictions of our LDDE model for ($\\Omega_m$,$\\Omega_\\lambda$)=(1.0,0.0)." + }, + "0112/astro-ph0112513_arXiv.txt": { + "abstract": "{Based on spectra obtained at the Anglo-Australian Observatory, we present a discussion of the metallicity of the galactic disc derived using Cepheids at galactocentric distances 4-6 kpc. Our new results together with previous gradient determination (Paper I) show that the overall abundance distribution within the galactocentric distances 4--11~kpc cannot by represented by a single gradient value. The distribution is more likely bimodal: it is flatter in the solar neighbourhood with a small gradient, and steepens towards the galactic center. The steepening begins at a distance of about 6.6 kpc. ", + "introduction": "In our previous work (Andrievsky et al. 2002, hereafter Paper I) results on elemental abundance distributions in the galactic disc based on 236 high-resolution spectra of 77 classical Cepheids in the solar neighborhood (galactocentric distances from 6 kpc to approximately 10.5 kpc) were reported. We found that among the 25 studied chemical elements, those from carbon to yttrium show small negative gradients, while heavier species produce near-to-zero gradients. Typical gradient values for iron-group elements were found to be equal to $\\approx -0.03$ dex~kpc$^{-1}$. In order to extend our previous study and to check the behavior of the elemental distribution towards the galactic center we have observed several Cepheids with galactocentric distances between 4-6 kpc. In this work we present the results from these stars and discuss them together with the data from Paper I. ", + "conclusions": "We supplemented our previous data on elemental abundance distributions in the solar neighborhood (Paper I) with new determinations based on Cepheids at distances of 4-6 kpc. Our new results together with previous gradient determinations (Paper I) show that the abundance distribution over the galactocentric distances 4--11~kpc cannot by represented by a single gradient value. More likely, the distribution is bimodal: it is flatter in the solar neighborhood with a small gradient, and becomes steeper towards the galactic center. The steepening begins at the distance about 6.5 kpc. \\newpage" + }, + "0112/astro-ph0112039_arXiv.txt": { + "abstract": "We study the nucleosynthesis and the induced mixing during the merging of massive stars inside a common envelope. The systems of interest are close binaries, initially consisting of a massive red supergiant and a main-sequence companion of a few solar masses. We apply parameterized results based on hydrodynamical simulations to model the stream-core interaction and the response of the star in a standard stellar-evolution code. Preliminary results are presented illustrating the possibility of unusual nucleosynthesis and post-merging dredge-up which can cause composition anomalies in the supergiant's envelope. ", + "introduction": "Common envelope (CE) evolution is one of the most interesting evolutionary stages of an initially close binary. It occurs when the two components of a binary system orbit inside an extended common envelope which is not in synchronous rotation with the embedded binary (Iben \\& Livio, 1993). A binary encounters a CE phase by one of three suggested channels: (1) a dynamical instability (Darwin instability) or a secular tidal instability (\\opencite{Hut}; \\opencite{Lai93}), (2) in nova systems because of the expansion of the nova shell engulfing the companion, and (3) as a consequence of dynamical mass transfer. Once a common envelope has formed, the secondary continues orbiting around the core of the primary within this common envelope. As the secondary is affected by various drag forces due to its interaction with the envelope, the orbit of the binary slowly shrinks. The transfer of angular momentum from the orbital motion to the envelope causes the spin-up of the envelope. The final result of this slow spiral-in depends on how much of the released orbital energy has been deposited into the envelope compared to the binding energy of the envelope. The deposited frictional energy may drive expansion of the envelope, causing its partial or complete ejection. If most of the envelope is ejected, the system survives as a close binary consisting of the core of the primary and the secondary. This evolutionary path provides the favourite channel for the formation of short-period binaries with compact components. Here we are primarily interested in the situation where the deposited energy is not sufficient to eject the common envelope. Then the spiral-in continues until the secondary starts to fill its own Roche lobe and begins to transfer mass to the core of the giant. Eventually, the binary will merge resulting in the formation of a rapidly rotating single star. During this merger, material from the secondary forms a stream which emanates from the secondary and falls towards the primary core, encountering an ambient medium with increasing pressure and density. This hydrogen-rich material may penetrate to a depth of about $10^{10}\\,$cm where the temperature of the ambient matter is as high as a few $10^8\\,$K. This deep penetration may have two major consequences: (a) the initiation of hydrogen burning through the hot-CNO cycle which may provide a neutron flux sufficient for efficient s-processing; (b) the injection/generation of high-entropy material near the primary core which may lead to the dredge-up of helium, eventually changing the surface composition of the merger product. In this contribution we present some preliminary results of the evolutionary calculations which involve the modelling of the merging phase. ", + "conclusions": "In our preliminary runs we have calculated two models, one with no entropy change in the stream and one with a high entropy change. The main differences between the two calculations can be explained in terms of how the hydrogen-rich material is injected into the hot, He-burning zone. In the case of no entropy change, the stream penetrates into the convective He-burning zone directly. This leads to an immediate very energetic response with very efficient nucleosynthesis. This causes a rapid expansion of the core of the primary and results in the widening of the secondary orbit and a temporary interruption of the mass transfer. If there is large generation of entropy in the stream, $k=0.4$, the stream cannot penetrate as deep and does not reach the convective He burning zone at the start of mass transfer. Since there is no dramatic expansion of the primary core, mass transfer continues steadily at an approximately constant rate. In the impact zone, a high-entropy region is generated. With time, a convective zone just below the impact zone starts to develop. Once created, this zone slowly erodes the core, mixing hydrogen down into hotter layers. The growth of this convective zone proceeds exponentially, very slow at the beginning, speeding up steadily up to the moment where it connects to the He-burning convective zone. At that point, hydrogen is quickly mixed throughout the He-burning zone. Unlike the previous case, the primary at this stage has accumulated much more hydrogen in this zone, which therefore leads to much more dramatic nuclear burning when the hydrogen is mixed into the helium-burning zone. There are many uncertainties which can affect and significantly change these results. The models discussed above have been calculated with the assumption that there is significant convective undershooting. Nevertheless, calculations which have been done for a smaller entropy change ($k=0.2$) but without undershooting have shown a behavior similar to the behaviour in the second model; but the model with $k=0.4$ and no undershooting did not show any mixing of hydrogen into the helium-burning zone, though the entropy generated in the core neighborhood during the merger can still induce mixing at a later stage. Another possible uncertainty is the speed of mixing in the convective zones, as it affects the rate with which the hydrogen-rich material is moving into the hot zone, which affects the nucleosynthesis. Our preliminary conclusion is that, during the merger, hydrogen-rich material from the secondary may be able to penetrate into the He-burning zone, though in the case of high-entropy generation in the stream, it is likely to happen only if there is significant undershooting." + }, + "0112/astro-ph0112569_arXiv.txt": { + "abstract": "This is the first in a series of papers on the effects of dust on the formation, propagation, and structure of nonlinear MHD waves and MHD shocks in weakly-ionized plasmas. We model the plasma as a system of 9 interacting fluids, consisting of the neutral gas, ions, electrons, and 6 grain fluids comprised of very small grains or PAHs and classical grains in different charge states. We formulate the governing equations for perpendicular shocks under approximations appropriate for dense molecular clouds, protostellar cores, and protoplanetary disks. We describe a code that obtains numerical solutions using a finite difference method, and establish its accuracy by comparing numerical and exact solutions for special cases. ", + "introduction": "Shock waves in weakly-ionized plasmas usually exhibit a multifluid structure, in which the charged and neutral particles behave as separate, interpenetrating fluids. If the shock speed is less than the speed at which the charged fluid communicates compressive disturbances, the charged particles will be accelerated ahead of the neutrals in a ``magnetic precursor,'' and the neutrals accelerated, compressed and heated downstream by collisions with streaming charged particles (Mullan 1971; Draine 1980, henceforth D80). The flow is continuous (a C- or C$^*$ shock) if the length scale for acceleration of the neutrals is large compared to the scale for cooling; otherwise the flow undergoes a discontinuous transition (a J shock) at a viscous subshock (D80; Chernoff 1987; Roberge \\& Draine 1990). The dynamics, chemistry, and emission from multifluid shock waves can differ profoundly from those of their single-fluid counterparts. Consequently, they have been the subject of numerous investigations (see Draine \\& McKee 1993 for the most recent comprehensive review). The energy and momentum of the ``charged fluid'' are usually dominated by the magnetic field, the charged particles acting mainly to provide energy and momentum transfer between the field and the neutral fluid. At low densities the coupling is dominated by ion-neutral scattering and shock waves can be modeled as a three-fluid system: the neutral gas, the ion/magnetic field fluid, and the fluid of electrons (which move with the ions but need not have the same temperature; see D80). However when the density of the preshock medium exceeds $\\nH \\sim 10^5$\\,\\cmMMM,\\footnote{We use \\nH\\ to denote the density of H nuclei in all forms.} the energy and momentum transfer is dominated by collisions between neutral particles and charged dust grains (Elmegreen 1979; D80; Nakano \\& Umebayashi 1980; Ciolek \\& Mouschovias 1993). If the grains are well coupled to the magnetic field, they can be incorporated into the 3-fluid picture as an additional species of ion. This is a reasonable approximation for PAHs and very small grains with radii $\\asg \\sim 10^{-7}$\\,cm (henceforth referred to collectively as ``small grains'') at densities up to $\\sim 10^9$ (Neufeld \\& Hollenbach 1994). If the small grains are present in dense gas with a fractional abundance not much less than $\\sim 10^{-7}$ (the abundance inferred for the diffuse ISM, see Li \\& Draine 2001 and references therein), they will dominate the field-neutral coupling and a 3-fluid model of the dynamics is appropriate (e.g., Kaufman \\& Neufeld 1996a,b). However if small grains are depleted in dense gas, as seems likely, then scattering by large grains with $\\ag \\simgt 0.1$\\,\\micron\\ will dominate the field-neutral coupling. Because the large grains are only partially tied to the magnetic field for $\\nH \\simgt 10^5$\\,\\cmMMM, they must be treated as an additional fluid (Chernoff 1985; Havnes, Hartquist \\& Pilipp 1987; Ciolek \\& Mouschovias 1993). The problem is further complicated by the fact that large grains have different charge states, a continuous spectrum of sizes, and masses many orders of magnitude larger than those of the ions. Draine (D80) developed an approximate description of the grain dynamics by treating the grains as test particles and neglecting their inertia; in this approximation each grain drifts through the neutrals with a speed that can be determined by balancing the gas drag and electromagnetic forces. Chernoff (1985) treated the grains as a separate fluid, included their inertia, and demonstrated that the grains can undergo a J shock. Pilipp, Hartquist \\& Havnes (1990) and Pilipp \\& Hartquist (1994) carried out numerical simulations of steady 4-fluid shocks, emphasizing the importance of grain charging processes and the dynamical effects of the grain charge and current. Wardle (1998, see also Wardle \\& Ng 1999) has formulated equations of motion valid for steady multifluid shocks with an arbitrary number of grain fluids, in a form that is valid for very large densities, and has discussed optimal procedures for representing the spectrum of large grain sizes by a single, effective size. This is the first in a series of papers on the effects of dust on the formation, propagation, and structure of nonlinear MHD waves and shocks in dense molecular clouds, bipolar outflows, and protoplanetary disks. Here we describe a dynamical model that incorporates and extends previous work in several ways: (i) We allow for time-dependent flow. Prior work on time-dependent, multifluid shocks (T\\'{o}th 1994; Smith \\& Mac Low 1997; Mac Low \\& Smith 1997; Stone 1997; Chieze, Pineau des Forets \\& Flower 1998) has been based on the 3-fluid model. We are interested in time-dependent solutions because they permit us to study the formation of magnetic precursors, because of their possible relevance to bipolar outflows (Flower \\& Pineau des Forets 1999), and because time-dependent simulations are a robust way to find steady solutions (e.g., steady J shocks) that are otherwise hard to compute. (ii) We include the inertia of large grains, which has been neglected in all prior studies except Chernoff's (1985). Neglecting the inertia is a good approximation when the gas drag time for large grains is small compared to all other dynamical time scales. While this is usually appropriate for the dynamical time of the neutrals, it is often a poor approximation for the charged fluid (see \\S2.3.3). A realistic treatment of grain dynamics is necessary for studies on the formation of magnetic precursors: Draine \\& McKee (1993) pointed out that magnetic precursors could not form (and C shocks therefore could not exist) for shocks faster than $\\approx 20~\\kms$ if large grains are able to ``load'' the magnetic field lines efficiently. The code we describe was designed in part to address this issue. (iii) We include small grains as a separate fluid. This permits us to model grain charging realistically and to examine various scenarios for small-grain depletion in dense gas. (iv) We include charge fluctuations by treating the grains in different charge states (3 each for small- and large grains) as separate fluids. Charge fluctuations exchange mass, momentum, and energy between the charged and neutral grain fluids, with implications for the loading of magnetic field lines. Nevertheless, the scope of our calculations is limited in some important respects: We do not attempt to describe the continuous size distribution of the large grains, which are assumed to have uniform radii. We have also restricted our initial investigations to one-dimensional geometry so that, for example, time-dependent simulations of Wardle instabilities with dust are beyond our present capabilities. The plan of this paper is as follows. In \\S~2 we give the equations of motion and discuss the conditions where various approximations are likely to be valid. In \\S~3 we describe our numerical methods and summarize a suite of benchmark calculations that establish the accuracy of our code. Our results are summarized in \\S~4. ", + "conclusions": "In this paper --- the first of a series --- we have discussed the formulation of a time-dependent, numerical MHD code designed to model the formation and evolution of perpendicular MHD shocks in dense molecular gas. Chief among the various features of our code is that it is a multifluid code, incorporating the dynamics of neutral, ion, and six different types of grain fluids. The grain fluids consist of populations with two distinct radii, each with three different charge states. In addition to including for the effects of elastic collisions with gas particles, we also account for the transfer of mass and momentum between the grain species due to capture of charged particles on grain surfaces. We have presented the equations required to model this physical system, and also described simplifying approximations adopted within the numerical code. To establish the accuracy of our numerical code, we have presented several test models and compared their results to exact, analytical solutions. In general, we found good agreement between the numerical and analytical solutions. This included a test of the formation of a C shock in a dust-free, weakly ionized gas, whose analytical solution was determined by Chernoff (1987). We also presented an illustrative model of a multifluid flow in which inertia of the large grains was shown to be important over a finite scale of the flow. The possible effects that differential motion of fluids in shocks may have on the state of the neutral gas and other species will be investigated in more detail elsewhere. The effect of grains (including their inertia) on the speed and propagation of signals in dense gas will also be a topic of discussion in future work. Using this numerical code as a tool to aid in further research, we intend in subsequent papers to study the relevant physics and chemistry of the formation and evolution of various types of nonlinear MHD waves and shocks in molecular clouds, protostellar cores, and disks." + }, + "0112/astro-ph0112043_arXiv.txt": { + "abstract": "We investigate the dependence of galaxy clustering on luminosity and spectral type using the 2dF Galaxy Redshift Survey (2dFGRS). Spectral types are assigned using the principal component analysis of Madgwick et al. We divide the sample into two broad spectral classes: galaxies with strong emission lines (`late-types'), and more quiescent galaxies (`early-types'). We measure the clustering in real space, free from any distortion of the clustering pattern due to peculiar velocities, for a series of volume-limited samples. The projected correlation functions of both spectral types are well described by a power law for transverse separations in the range 2$<$($\\sigma/h^{-1}$\\,Mpc)$<$15, with a marginally steeper slope for early-types than late-types. Both early and late types have approximately the same dependence of clustering strength on luminosity, with the clustering amplitude increasing by a factor of $\\sim$2.5 between $L^*$ and 4$L^*$. At all luminosities, however, the correlation function amplitude for the early-types is $\\sim$50\\% higher than that of the late-types. These results support the view that luminosity, and not type, is the dominant factor in determining how the clustering strength of the whole galaxy population varies with luminosity. ", + "introduction": "\\label{sec:intro} One of the major goals of large redshift surveys like the 2 degree Field Galaxy Redshift Survey (2dFGRS) is to make an accurate measurement of the spatial distribution of galaxies. The unprecedented size of the 2dFGRS makes it possible to quantify how the clustering signal depends on intrinsic galaxy properties, such as luminosity or star formation rate. The motivation behind such a program is to characterize the galaxy population and to provide constraints upon theoretical models of structure formation. In the current paradigm, galaxies form in dark matter haloes that are built up in a hierarchical way through mergers or by the accretion of smaller objects. The clustering pattern of galaxies is therefore determined by two processes: the spatial distribution of dark matter haloes and the manner in which dark matter haloes are populated by galaxies (Benson \\etal 2000b; Peacock \\& Smith 2000; Seljak 2000; Berlind \\& Weinberg 2002). The evolution of clumping in the dark matter has been studied extensively using N-body simulations of the growth of density fluctuations via gravitational instability (e.g. Jenkins \\etal 1998; 2001). With the development of powerful theoretical tools that can follow the formation and evolution of galaxies in the hierarchical scenario, the issue of how galaxies are apportioned amongst dark matter haloes can be addressed, and detailed predictions of the clustering of galaxies are now possible (Kauffmann, Nusser \\& Steinmetz 1997; Kauffmann \\etal 1999; Benson \\etal 2000a,b; Somerville \\etal 2001). The first attempt to quantify the difference between the clustering of early and late-type galaxies was made using a shallow angular survey, the Uppsala catalogue, with morphological types assigned from visual examination of the photographic plates (Davis \\& Geller 1976). Elliptical galaxies were found to have a higher amplitude angular correlation function than spiral galaxies. In addition, the slope of the correlation function of ellipticals was steeper than that of spiral galaxies at small angular separations. More recently, the comparison of clustering for different types has been extended to three dimensions using redshift surveys. Again, similar conclusions have been reached in these studies, namely that ellipticals have a stronger clustering amplitude than spirals (Lahav \\& Saslaw 1992; Santiago \\& Strauss 1992; Iovino \\etal 1993; Hermit \\etal 1996; Loveday \\etal 1995; Guzzo \\etal 1997; Willmer \\etal 1998). The subjective process of visual classification can now be superseded by objective, automated algorithms to quantify the shape of a galaxy. One recent example of such a scheme can be found in Zehavi \\etal (2002), who measured a ``concentration parameter'' for $30\\,000$ galaxy images from the Sloan Digital Sky Survey, derived from the radii of different isophotes. Again, based upon cuts in the distribution of concentration parameter, early-types are found to be more clustered than late-types. In this paper, we employ a different method to classify galaxies, based upon a principal component analysis (PCA) of galaxy spectra, which is better suited to the 2dFGRS data (Madgwick \\etal 2002). This technique has a number of attractive features. First, the PCA approach is completely objective and reproducible. An equivalent analysis can, for example, be applied readily to spectra produced by theoretical models of galaxy formation or to spectra obtained in an independent redshift survey. Secondly, the PCA can be applied over the full magnitude range of the 2dFGRS, whenever the spectra are of sufficient signal to noise (see Section~\\ref{subsec:eta_class}). For the 2dFGRS, the image quality is adequate to permit a visual determination of morphological type only for galaxies brighter than $b_{\\rm J}\\,\\simeq\\,17$, which comprise a mere $5\\%$ of the spectroscopic sample. Two previous clustering studies have used spectral information to select galaxy samples. Loveday, Tresse \\& Maddox (1999) grouped galaxies in the Stromlo-APM redshift survey into three classes based upon the equivalent width of either the \\Halpha\\ or OII lines, and found that galaxies with prominent emission lines display weaker clustering than more quiescent galaxies. Tegmark \\& Bromley (1999) measured the relative bias between different spectral classes in the Las Campanas redshift survey (Shectman \\etal 1996), using a classification derived from PCA analysis (Bromley \\etal 1998), and also found that early spectral types are more strongly clustered than late spectral types. (See also Blanton 2000 for a revision of Tegmark \\& Bromley's analysis, which takes into account the effect of errors in the survey selection function.) Here, we use the 2dFGRS survey to measure the dependence of galaxy clustering jointly on luminosity and spectral type, adding an extra dimension to the analysis carried out by \\paperone. Previously, a pioneering study of bivariate galaxy clustering, in terms of luminosity and morphological type, was carried out using the Stromlo-APM redshift survey (Loveday \\etal 1995). To place the analysis presented here in context, the samples that we consider cover a larger volume and, despite being volume-limited (see Section~\\ref{sec:vol.lim}), typically contain over an order of magnitude more galaxies than those available to Loveday \\etal We give a brief overview of the 2dFGRS in Section~\\ref{sec:data}, along with details of the spectral classification and an explanation of how the samples used in the clustering analysis were constructed. The estimation of the redshift space correlation function and its real space counterpart, the projected correlation function, are outlined in Section~\\ref{sec:meth}. A brief overview of the clustering of 2dFGRS galaxies in redshift space, selected by luminosity and spectral type, is given in Section~\\ref{sec:redspace}; a more detailed analysis of the redshift space clustering can be found in Hawkins \\etal (2002). We present the main results of the paper in Section~\\ref{sec:res} and conclude in Section~\\ref{sec:conc}. ", + "conclusions": "\\label{sec:conc} \\begin{figure} \\epsfig{file=class.mix.norm.ps,width=0.47\\textwidth,clip=,bbllx=50,bblly=205,bburx=550,bbury=640} \\caption{ The fraction of galaxies in the two broad spectral classes, early-type and late-type, as a function of absolute magnitude. The fractions are derived from the volume-limited samples listed in Tables 1, 2 and 3. The error bars show the Poisson errors on the fractions. } \\label{fig:ndens} \\end{figure} We have used the 2dFGRS to study the dependence of clustering on spectral type for samples spanning a factor of twenty in galaxy luminosity. The only previous attempt at a bivariate luminosity-morphology/spectral type analysis of galaxy clustering was performed by Loveday \\etal (1995) using the Stromlo-APM redshift survey. They were able to probe only a relatively narrow range in luminosity around $L^{\\star}$, which is more readily apparent if one considers the median magnitude of each of their magnitude bins (see Fig. 3b of Norberg \\etal 2001). The scatter between spectral and morphological types illustrated in Fig.~\\ref{fig:jon.eta} precludes a more detailed comparison of our results with those of earlier studies, based on morphological classifications. In Norberg \\etal (2001), we used the 2dFGRS to make a precise measurement of the dependence of galaxy clustering on luminosity. The clustering amplitude was found to scale linearly with luminosity. One of the aims of the present paper is to identify the phenomena that drive this relation. In particular, there are two distinct hypotheses that we wish to test. The first is that there is a general trend for clustering strength to increase with luminosity, regardless of the spectral type of the galaxy. The second is that different types of galaxies have different clustering strengths, which may vary relatively little with luminosity, but a variation of the mix of galaxy types with luminosity results in a dependence of the clustering strength on luminosity. Madgwick \\etal (2002) estimated the luminosity function of 2dFGRS galaxies for different spectral classes, and found that, in going from early-type to late-types, the slope of the faint end of the luminosity function becomes steeper while the characteristic luminosity becomes fainter. Another representation of the variation of the luminosity function with spectral class is shown in Fig.~\\ref{fig:ndens}, where we plot the fraction of early and late type galaxies in absolute magnitude bins. The plotted fractions are derived from the volume-limited samples listed in Tables~1,~2 and~3. The mix of spectral types changes dramatically with luminosity; faint samples are dominated by late-types, whereas early-types are the most common galaxies in bright samples. Similar trends were found for galaxies labelled by morphological type in the SSRS2 survey by Marzke \\etal (1998). \\begin{figure} \\epsfig{file=bias.ps,width=0.47\\textwidth,clip=,bbllx=47,bblly=225,bburx=560,bbury=672} \\caption{ The relative bias (on the scale of $r = 4.89$\\mpc) of the different spectral classes as a function of luminosity, as indicated by the key. The definition of relative bias is given in the text. The reference sample covers the absolute magnitude range $-19.5\\,\\ge\\,M_{b_{\\rm J}}-5\\log_{10}\\,h\\,\\ge\\,-20.5$. The fiducial luminosity, $L^{\\star}$, is taken to be $M_{b_{\\rm J}}-5\\log_{10}\\,h\\,=-19.7$. The solid line shows the fit to the results of \\paperone, $b/b^{\\star} = 0.85 + 0.15L/L^{\\star}$, whereas the dashed line shows the fit to the data analyzed here. All the error bars plotted take into account the correlation between the various samples. } \\label{fig:bl} \\end{figure} We find that the change in the mix of spectral types with luminosity is not the main cause for the increase in the clustering strength of the full sample with luminosity. To support this assertion, we plot in Fig.~\\ref{fig:bl} the variation of clustering strength with luminosity normalized, for {\\it each} spectral class, to the clustering strength of a fiducial sample of \\Mstar\\ galaxies, {\\it i.e.} the sample which covers the magnitude range $-19.5\\,\\ge\\,M_{b_{\\rm J}}-5\\log_{10}\\,h\\,\\ge\\,-20.5$. For a galaxy sample with best fitting correlation function parameters $r^{i}_{0}$ and $\\gamma^{i}$, we define the relative bias with respect to the \\Mstar\\ sample of the same type by \\begin{equation} \\frac{b^{i}}{b^{\\star}}\\Big|_{\\eta}(r) = \\sqrt{\\frac{ (r^{i}_{0})^{\\gamma_{i}}}{r^{\\gamma}_{0}} r^{\\gamma-\\gamma_{i}}}\\Big|_{\\eta}, \\label{eq:bias} \\end{equation} where $r_{0}$ and $\\gamma$ are the best fitting power-law parameters for the fiducial sample. In Fig.~\\ref{fig:bl}, we plot the relative bias evaluated at a fixed scale, $r = 4.89$\\mpc, which is the correlation length of the reference sample for all \\etapar-classified galaxies. A scale dependence in Eq.~\\ref{eq:bias} arises if the slopes of the real space correlation functions are different for the galaxy samples being compared. In practice, the term $r^{\\gamma-\\gamma_{i}}$ is close to unity for the samples considered. The dashed line shows a fit to the bias relation defined by the open symbols. The solid line shows the effective bias relation obtained by \\paperone, which is defined in a slightly different way to the effective bias computed here. From Fig.~\\ref{fig:bl}, we see that the trend of increasing clustering strength with luminosity in both spectral classes is very similar for galaxies brighter than $L>0.5L^{\\star}$. At the brightest luminosity, corresponding to $\\sim\\,4 L^{\\star}$, the clustering amplitude is a factor of $2-2.5$ times greater than at $L^{\\star}$. This increase is much larger than the offset in the relative bias factors of early and late types at any given luminosity. We conclude that the change in correlation length with absolute magnitude found by Norberg \\etal (2001) is primarily a luminosity effect rather than a reflection of the change in the mix of spectral types with luminosity. Benson \\etal (2001) showed that a dependence of clustering strength on luminosity is expected in hierarchical clustering cold dark matter universes because of the preferential formation of the brightest galaxies in the most massive, strongly clustered dark halos. The close connection between the spectral characteristics of galaxies and their clustering properties discussed in this paper provides further evidence that the galaxy type is also related to the mass of the halo in which galaxy forms." + }, + "0112/astro-ph0112333_arXiv.txt": { + "abstract": "{Long-lasting, intense, stellar X--ray flares may approach conditions of breaking magnetic confinement and evolving in open space. In the perspective of searching for possible tracers of non-confinement, we explore this hypothesis with hydrodynamic simulations of flares occurring in a non-confined corona: model flares are triggered by a transient impulsive heating injected in a plane-parallel stratified corona. The plasma evolution is described by means of a numerical 2-D model in cylindrical geometry $R,Z$. We explore the space of fundamental parameters. As a reference model, we consider a flare triggered by a heating pulse of 10 erg cm$^{-3}$ s$^{-1}$ lasting 150 s and released in a region $\\sim 10^9$ cm wide and at a height $\\sim 2 \\times 10^9$ cm from the base of the stellar surface. The pressure at the base of the corona of the unperturbed atmosphere is 0.1 dyne cm$^{-2}$. The heating would cause a 20 MK flare if delivered in a 40000 km long closed loop. The modeled plasma evolution in the heating phase involves the propagation of a 10 MK conduction front and the evaporation of a shocked bow density front upwards from the chromosphere. As the heating is switched off, the temperature drops in few seconds while the density front still propagates, expanding, and gradually weakening. This kind of evolution is shared by other simulations with different coronal initial pressure, and location, duration and intensity of the heating. The X-ray emission, spectra and light curves at the ASCA/SIS focal plan, and in two intense X--ray lines (Mg XI at 9.169~\\AA~and Fe XXI at 128.752~\\AA), have been synthesized from the models. The results are discussed and compared to features of confined events, and scaling laws are derived. The light curves invariably show a very rapid rise, a constant phase as long as the constant heating is on, and then a very fast decay, on time scales of few seconds, followed by a more gradual one (few minutes). We show that this evolution of the emission, and especially the fast decay, together with other potentially observable effects, are intrinsic to the assumption of non--confinement. Their lack indicates that observed long--lasting stellar X--ray flares should involve plasma strongly confined by magnetic fields. ", + "introduction": "\\label{sec:intro} Stellar coronae are not spatially resolved with present day observations. It is widely accepted that they share many of the basic mechanisms of the solar corona, the only resolved corona that we know. Decades of X-ray observations have shown that both its structuring and heating are dominated by the ambient magnetic field. Evidence has been found, such as the presence of coronal plasma at temperatures of several million degrees (e.g. Linsky 1981, Schmitt et al. 1990), that some form of confinement of X-ray bright plasma should be at work on stars, probably associated with intense magnetic fields. It is then natural to wonder what is the morphology of the confining structures in a stellar corona. On the Sun, plasma is typically confined in magnetic arches (the loops), and both the geometry and luminosity of many of them are steady for relatively long times. If stellar coronae are similar to the solar one, then structures similar to solar loops should be found on stars. Several approaches are possible to determine the dimensions of coronal structures. An eclipse on Algol has allowed to find the volume responsible of a flare (Schmitt \\& Favata 1999). Loop fitting has been used to constrain the average linear dimensions of stellar coronal loops (e.g. Maggio \\& Peres 1997). Another source of information on the characteristic of stellar coronal loops is the evolution of X-ray flares. In particular, basic mechanisms of plasma cooling are radiation and thermal conduction from the corona downwards to the cooler chromosphere. Linking the plasma characteristic cooling times to the observed flare decay time allows to infer the linear size of the cooling region. This property has been largely exploited in the past to estimate the size of stellar flaring loops (e.g. Haisch 1983, van den Oord et al. 1988, van den Oord \\& Mewe 1989, Reale et al. 1988, Reale et al. 1997, Favata \\& Schmitt 1999, see Reale 2001 for a review). One implication of the simple scaling from cooling times is that slower and slower decays, associated with long-lasting flares, correspond to longer and longer loops. As a matter of fact, stellar X-ray flares are typically observed to last much longer than average solar X-ray flares, up to several days whereas the duration of solar flares ranges from few minutes to several hours. In spite of the obvious observational bias toward detecting long events, the evidence of long flares has been taken as strong indication of very large flaring regions, in particular comparable to, or even larger than, the stellar radius (e.g. Graffagnino et al. 1995, Favata \\& Schmitt 1999, Tsuboi et al. 2000, Favata 2001). On the Sun, large scale and long-duration events typically occur in complex coronal structures, in which the magnetic field undergoes major rearrangements and several loop structures are progressively involved (e.g. two-ribbon flares): the large amounts of energy released lead to large plasma pressure which may exceed the magnetic one and cause the end of confinement. The question now is: may plasma in long-lasting and intense stellar flares break the magnetic cage, and erupting, be free to move and evolve in open space? This work addresses this question, by modeling flares occurring in a totally non-confined stratified corona. Although we do not pretend this to be an entirely realistic scenario, it can be considered as an extreme limit of situations in which the plasma dynamics dominates over the confining magnetic field and therefore governs the X-ray luminosity evolution. Indeed, some solar and stellar flares show phenomena that indicate that at least some of the plasma is not confined. In the solar case, coronal mass ejections (CMEs) following flares are not confined. In the stellar case VLBI measurements of nonthermal particles propagating away from the star are observed. A good example is a flare in UX Ari observed by Mutel et al (1985). White and Franciosini (1995) also show that the relativistic plasma responsible for the gyrosynchrotron radio emission in stellar comes from expanding plasma. Here we focus on studying the X-ray emitting plasma. Our scope is to explore: to what point actual stellar flares can approximate free-expanding flaring plasma, if there are distinctive signatures in X-ray emission (obtained from the modeling) which characterize unconfined flares vs confined ones and therefore allow us to discriminate the confining role of the magnetic field, and if observed stellar flares present such signatures. The study that we illustrate here presents also interesting results on the theoretical point of view: we describe the evolution of the non-confined plasma in a stratified coronal atmosphere, i.e. which is hot {\\it ab initio} and where thermal conduction is fully efficient. This aspect makes this work quite different from typical models of bursts, like supernovae, which instead propagate in non-conducting and cooler media. Our approach is to set up an initial stratified atmosphere, similar to a confined one, including chromosphere, transition region and corona, and where the plasma is not confined to move along one spatial direction, and then to impart a localized energy impulse. Plasma dynamics and heat conduction can occur isotropically. The evolution of such system needs to be described by a numerical 2-D hydrodynamic model, with enhanced spatial resolution in the low corona and in the transition region and including plasma thermal conduction; therefore it is also a demanding numerical problem. In order to compare our results with observations and to obtain diagnostical feedbacks, we synthesize from our model results the expected evolution of the integrated X-ray flare luminosity, i.e. by simulating a flaring unresolved source, such as a star. To have a realistic scenario, representative of real flare observations, we have chosen to synthesize the emission in a wide X-ray band such as at the focal plane of the ASCA/SIS (sensitive in the range, approximately, 0.3 - 10 keV), and in two intense X-ray lines. Similar results can be expected for other wide band CCD instruments such as Chandra/ACIS and XMM/EPIC, while the lines may be detectable by instruments with good spectral resolution such as Chandra and XMM-Newton grating detectors. In Sect.~\\ref{sec:model} we describe our modeling approach in detail, in Sect.~\\ref{sec:simul}, the simulations performed and the results obtained are presented, in Sect.~\\ref{sec:discuss} the results are discussed and in Sect.~\\ref{sec:conclu} conclusions are drawn. ", + "conclusions": "\\label{sec:conclu} Stellar flares are generally observed to evolve and decay on time scales ranging from several hours to days, suggesting large flaring regions or even lack of magnetic confinement. This work shows that the hydrodynamic evolution of flares occurring in non-confined atmospheres leads invariably to a much faster decay of the brightness (on time scales of very few minutes) after the heating phase, in a wide range of the physical parameters. One of the main implications is that the long duration of stellar flares is not indicative of non-confined plasma. Reversing the argument, it is highly probable that the observed long-lasting stellar X-ray flares involve forms of plasma confinement in closed coronal structures. We have pointed out other characteristic features, with more or less diagnostical power, e.g. the synchronous evolution of density and temperature, summarized in Table~\\ref{tab:disc}, that seem to indicate that X-ray observed stellar flares involve mostly plasma confined in closed structures. Although this work simulates rather extreme conditions of flares occurring in completely non-confined atmospheres, the result that long-lasting stellar flares are likely not occurring in large open structures largely motivates its validity. Furthermore this work predicts the possible existence of a new phenomenological class of events, characterized by peculiar light curves with a limited range of time scales. Such kind of events may be addressed by observations at high sensitivity levels, which allow for high time resolution and high signal-to-noise ratio, such as those obtained from the current missions Chandra and XMM-Newton. \\bigskip \\bigskip" + }, + "0112/astro-ph0112356_arXiv.txt": { + "abstract": "In this paper we analyze the relation between radio, optical continuum and {\\HalphaNII} emission from the cores of a sample of 21 nearby Fanaroff \\& Riley type I galaxies as observed with the VLBA and HST. The emission arises inside the inner tens of parsec of the galaxies. Core radio emission is observed in 19/20 galaxies, optical core continuum emission is detected in 12/21 galaxies and {\\HalphaNII} core emission is detected in 20/21 galaxies. We confirm the recently detected linear correlation between radio and optical core emission in FR I galaxies and show that both core emissions also correlate with central {\\HalphaNII} emission. The tight correlations between radio, optical and {\\HalphaNII} core emission constrain the bulk Lorentz factor to $\\gamma \\sim 2 -5$ and $\\gamma \\lta 2$ for a continuous jet and a jet consisting of discrete blobs, respectively, assuming jet viewing angles in the range $[30\\deg,90\\deg]$. Radio and optical core emissions are likely to be synchrotron radiation from the inner jet, possibly with a significant contribution from emission by an accretion disk and/or flow. Elliptical galaxies with LINER nuclei without large-scale radio jets seem to follow the core emission correlations found in FR I galaxies. This suggests that the central engines could be very similar for the two classes of AGNs. ", + "introduction": "\\label{s:intro} Many galaxies in the nearby universe ($z<0.1$) contain an active galactic nucleus (AGN) which features a central low-ionization narrow emission-line region (LINER). These AGNs typically have a bolometric luminosity $L_{\\rm bol} \\lta 10^{43} \\ergs$ which is orders of magnitude smaller than those of classical AGN at higher redshifts, such as Seyferts, Quasars and Broad-line Radio Galaxies. The energy for the activity is thought to be provided by accretion of matter onto a supermassive black hole residing at the galaxy center in all AGN classes. Radio galaxies constitute a subset of the active galaxy family. These galaxies have radio cores at their centers from which jets emerge. Radio galaxies are divided in two classes based on differences in large-scale radio morphology: low-power Fanaroff \\& Riley (1974) type I (FR I) radio galaxies and high-power radio FR II galaxies (cf.~Bridle \\& Perley 1984). Many nearby FR I nuclei show optical LINER emission (e.g., Baum, Zirbel \\& O'Dea 1995). The goal of this paper is to study the active cores of nearby FR I galaxies and the relation to their nearby LINER-type counterparts without radio jets. The hosts of FR I galaxies are almost without exception bright early-type galaxies (see Ledlow \\etal 2001 for counter example). Quiescent and FR I radio galaxies in nearby Abell clusters show no statistical differences in their global properties (Ledlow \\& Owen 1995). The two classes do not differ in surface-brightness profiles, the surface-brightness and size relations, the distribution over ellipticities, and the occurrence and strength of non-elliptical isophotes. Zirbel (1993, 1996) found that FR I galaxies are larger than radio-quiet galaxies at given isophotal magnitude, but this result was not confirmed by Ledlow \\& Owen (1995). Ledlow \\& Owen (1995) did not find differences between the radio-loud and quiescent cluster galaxies in the local density of nearby companions, and in the frequency of morphological peculiarities or tidal interactions . Studies of samples of FR I galaxies either in the field or in small groups found indications of tidal interactions in a majority of them (Gonzalez-Serrano \\& Carballo 1993; Colina \\& de Juan 1995). It is unknown if this rate is different for quiescent galaxies in similar environments. In summary, it is unlikely that host environment and host global properties are decisive for the capability of an early-type galaxy to become a FR I galaxy, but tidal interactions might trigger the nuclear activity, at least in some cases. With the advent of HST it has become possible to look for host galaxy differences closer to the AGN itself. In a previous paper (Verdoes Kleijn \\etal 1999) we presented WFPC2 broad- and narrow-band images of a well-defined complete sample of 21 nearby FR I early-type UGC galaxies, the `UGC FR I sample'. Similar WFPC2 surveys of radio galaxies have been carried out also for the 3CR sample (Martel \\etal 1999; de Koff \\etal 2000) and the B2 sample (Capetti~\\etal 2000). The UGC FR I sample WFPC2 data suggest that the stellar content (i.e., color and shape) is very similar to that of quiescent early-type galaxies including the central few kiloparsec. However, a rigorous comparison with a quiescent galaxy sample to discern subtle differences has not yet been performed. Potential AGN fuel is clearly present in the inner hundreds of parsecs. Central {\\HalphaNII} emission is detected in all sample galaxies and dust is detected in 19/21 galaxies. Other samples of FR I galaxies show emission-line detection rates of $\\gta 80\\%$ (e.g., Baum \\& Heckman 1989; Morganti \\etal 1992). The detection rate of small-scale dust in quiescent early-type galaxies is typically only $\\sim 40\\%$ (van Dokkum \\& Franx 1995; Tran~\\etal 2001), while emission-line gas is detected in typically $\\sim 60\\%$ of the quiescent early-type galaxy population (Philips \\etal 1986; Goudfrooij~\\etal 1993). Thus, the detection rate of potential fuel for an AGN is lower but significant in quiescent galaxies and hence its presence seems not to be the only factor determining the on-set of activity. A VLBA study of the radio properties of the cores and the inner jets in the UGC FR I nuclei has been presented by Xu \\etal (2000). In this paper we analyze the correlation between this AGN radio emission and optical continuum and {\\HalphaNII} core emission as observed with HST/WFPC2. We compare the UGC FR I core properties to those of 3CR FR I cores which have been studied in the radio, optical and X-rays (e.g., Chiaberge, Capetti \\& Celotti 1999; Hardcastle \\& Worrall 2000). For the 3CR sample, the correlations between the core emissions have led to favor the inner jet as their origin. We determine the validity for different AGN components to produce the core emissions in the UGC FR I galaxies. We also compare the FR I core properties to those of nearby LINER-type AGNs without large-scale radio jets. The outline of this paper is as follows. Section~\\ref{s:data} describes the sample and data from the literature used in this analysis. It also describes in detail the optical core flux measurements. Section~\\ref{s:correlations} quantifies the correlations between the core fluxes and shows they are not significantly affected by the ubiquitous central dust distributions. Section~\\ref{s:origincore} describes the arguments against the existence of compact stellar clusters producing the optical core emission. We present evidence that the radio and optical cores are synchrotron emission from the inner jet and discuss constraints on its Doppler boosting. The contribution to the core emission by an accretion disk and/or flow is also considered. Section~\\ref{s:halphanii} discusses the possible excitation mechanisms for the central gas. Section~\\ref{s:fr1sandllagns} surveys the similarities between the FR I and other LINER cores. Section~\\ref{s:summary} summarizes the main results and final conclusions. Appendix~\\ref{a:isophotes} presents the WFPC2 imaging analysis for UGC 7115 and UGC 12064 (3C449). Appendix~\\ref{a:beaming} discusses details of the beaming models used in Section~\\ref{s:boosting}. Throughout the paper we use a Hubble constant $H_0$= 75 kms$^{-1}$Mpc$^{-1}$. ", + "conclusions": "\\label{s:summary} In this paper we have analyzed the relation between $0.01''$ scale radio and $0.1''$ scale optical continuum and {\\HalphaNII} core emission of a complete sample of 21 nearby FR I galaxies. The main conclusions are: \\begin{enumerate} \\item{We confirm the linear correlation between optical and radio core emission in nearby FR I nuclei. We find that both core emissions also correlate with the core {\\HalphaNII} emission. The mutual correlations are unlikely to be caused by obscuration from the ubiquitous central dust.} \\item{ Nuclear stellar clusters are highly unlikely to be the source of the optical core emission for two reasons. First, previous spectral studies have directly excluded a nuclear stellar cluster for two typical members of the UGC FR I sample. Second, the UGC FR I radio, optical and {\\HalphaNII} core luminosities resemble the nuclei of AGN-type LINERS more closely than stellar-type LINERS. An AGN origin for the optical cores in all UGC FR I galaxies is strongly suggested by the tight correlation with radio core emission, which is certainly produced by the AGN.} \\item{ A jet origin for both the radio and optical core emission is favored because (i) optical and radio core emission are tightly correlated, (ii) spectral indices from radio to optical are similar to those for extended optical jets, (iii) there is a suggestive trend with independent estimates from jet orientation, and (iv) the correlation residuals from both core emission with {\\HalphaNII} emission are corrrelated. However, a significant contribution from a second component, such as accretion disk/flow/wind, to the radio and optical emission at the $0.1''$ scale for the optical and the $0.01''$ scale for the radio emission might be present.} \\item{The correlation of the optical and radio core emission with the isotropic {\\HalphaNII} emission constrains the core bulk Lorentz factor $\\gamma \\lta 2$ if the inner jets consist of discrete blobs ($p=3$). For a continuous jet, bulk Lorentz factors $\\gamma \\sim 2-5$ are inferred. This result crtically depends on the assumed range in viewing angle, which is assumed to be $[30\\deg,90\\deg]$. The bulk Lorentz factors required by the BL Lac - FR I unification scheme are generally larger, i.e..~$\\gamma > 5$. If the core emission is dominated by a jet component, this discrepancy could be reconciled by a two-layer jet with a fast moving spine surrounded by a slower outer layer.} \\item{The central gas is excited by AGN-related processes. Both shock- and photo-ionization appear plausible excitation mechanisms at this time.} \\item{Radio, optical and {\\HalphaNII} core luminosities of elliptical LINER-type AGNs with and without kiloparsec-scale radio jets appear to have similar relations. The results suggest (i) the engines in the two types producing the cores might be similar and (ii) the core radio emission is dominated by inner jet emission.} \\end{enumerate} Conclusions 1 to 5 prompt us to make the following speculation: if the radio and optical core emission are indeed inner jet synchrotron emission, then their strong correlation with {\\HalphaNII} core emission implies either a direct link between jet luminosity and gas excitation power, possibly via jet - gas interactions or a close relation between AGN photo-ionizing power and jet radiative (and possibly kinetic) power. Conclusion 6 and the close resemblance of quiescent and radio-loud active galaxies from scales just outside the AGN to the entire galaxy and even environment lead to two further speculations. First, either all bright early-type galaxies can host an AGN, or the capability of a bright early-type galaxy to become an AGN is set by conditions within the central few parsecs, currently unresolved at most wavelengths. Second, the capability of a LINER-type AGN to become radio-loud is set by conditions at similar scales. Two main ingredients for an active nucleus are present at these scales: accreting matter and a supermassive black hole . Likely factors determining the formation of large-scale radio jets are then the black hole spin and/or the inner accretion disk properties. A rigorous comparison between the nuclei of active and quiescent nearby galaxies at the tens of parsec resolution could advance our understanding on these issues." + }, + "0112/astro-ph0112483_arXiv.txt": { + "abstract": "The luminous X-ray binary Circinus~X-1 has been observed twice near zero orbital phase using the High-Energy Transmission Grating Spectrometer (HETGS) onboard {\\sl Chandra}. The source was in a high-flux state during a flare for the first observation, and it was in a low-flux state during a dip for the second. Spectra from both flux states show clear P~Cygni lines, predominantly from H-like and He-like ion species. These indicate the presence of a high-velocity outflow from the Cir~X-1 system which we interpret as an equatorial accretion-disk wind, and from the blueshifted resonance absorption lines we determine outflow velocities of 200--1900~km~s$^{-1}$ with no clear velocity differences between the two flux states. The line strengths and profiles, however, are strongly variable both between the two observations as well as within the individual observations. We characterize this variability and suggest that it is due to both changes in the amount of absorbing material along the line of sight as well as changes in the ionization level of the wind. We also refine constraints on the accretion-disk wind model using improved plasma diagnostics such as the He-like Mg~XI triplet, and we consider the possibility that the X-ray absorption features seen from superluminal jet sources can generally be explained via high-velocity outflows. ", + "introduction": "The discovery of broad P~Cygni line profiles from the luminous X-ray binary Circinus~X-1 (hereafter Cir~X-1), which we reported in Brandt \\& Schulz (2000; hereafter Paper~I), was the first time such lines had been observed clearly in the X-ray spectrum of a cosmic object. The P~Cygni lines demonstrated the presence of a high-velocity outflow, in line with other suggestions of outflow in this system (e.g., Johnston, Fender, \\& Wu 1999). However, the nature of Cir~X-1 in general is still poorly understood and, despite advances in recent years, there remains uncertainty about even the most basic properties of this system. Since its discovery (Margon et~al. 1971), it has appeared bright and variable in X-rays exhibiting a period of 16.6 days (Kaluzienski et~al. 1976), and it has thus been a frequent target of most X-ray observatories. The compact object in the Cir~X-1 system is thought to be a neutron star (Tennant, Fabian, \\& Shafer 1986) that can radiate at super-Eddington luminosities during times of strong mass transfer. Its heavily reddened optical counterpart (e.g., Moneti 1992) shows strong, asymmetric H$\\alpha$ emission; this emission is variable and appears to arise from several sites in the system including the accretion disk (e.g., Whelan et~al. 1977; Mignani, Caraveo, \\& Bignami 1997; Johnston et~al. 2001). The system shows two arcminute-scale radio jets (Stewart et~al. 1993), and an arcsecond-scale asymmetric jet (Fender et~al. 1998) suggests the presence of relativistic outflow from the source. Cir~X-1 is often included among the ``Galactic microquasar'' X-ray binaries (Mirabel 2001). The nature of the companion star is an important issue for the interpretation of the observed X-ray P~Cygni lines. Specifically, one must address if the P~Cygni lines could be made by absorption of X-rays from the neutron star in a high-velocity wind from the companion; such a scenario is likely for the high-mass X-ray binary Cygnus~X-3, which also shows X-ray P~Cygni lines (Liedahl et~al. 2000) and has a Wolf-Rayet star companion. The large velocities observed for the X-ray P~Cygni lines of Cir~X-1 (up to $\\pm 1900$~km~s$^{-1}$) would require a high-mass companion, probably of spectral type~O. Several attempts have been made to determine the spectral type of the companion star. While a low-mass companion is favored by most of the recent studies (e.g., Stewart et~al. 1991; Glass 1994), the constraints remain weak due largely to the heavy interstellar reddening. Additional evidence for a low-mass companion comes from observations of the correlated X-ray spectral and timing properties, observations of type~I X-ray bursts, and theoretical considerations. For example, Shirey, Bradt, \\& Levine (1999a) and Shirey, Levine, \\& Bradt (1999b) have demonstrated in extensive studies of \\xte data that Cir~X-1 exhibits spectral branches in the hardness-intensity diagram that can be identified with the horizontal, normal, and flaring branches of ``Z'' type low-mass X-ray binaries (LMXBs; e.g., Hasinger \\& van der Klis 1989; Schulz, Hasinger, \\& Tr\\\"umper 1989). Qu, Yu, \\& Li (2001) have also noted some similarities to the LMXBs GX~$5-1$ and Cyg~X-2. Furthermore, Tauris et~al. (1999) considered the (somewhat uncertain) kinematic properties of the Cir~X-1 system and argued that the companion is a low-mass ($\\lesssim 2.0$~$M_\\odot$), unevolved star. Given that the stellar wind from the likely low-mass companion cannot explain the observed properties of the X-ray P~Cygni profiles from Cir~X-1, we interpreted these profiles in the context of an equatorial wind driven from the accretion disk by a combination of Compton heating and radiation pressure (see Paper~I). The fact that Cir~X-1 can radiate with high luminosity relative to its Eddington luminosity ($L/L_{\\rm Edd}$) makes it a natural system in which to expect observable outflows, because of the larger amount of photon pressure available per unit gravitational mass. Accretion-disk winds in X-ray binaries have been discussed both theoretically and observationally (e.g., Begelman, McKee, \\& Shields 1983; Raymond 1993; Chiang 2001; Proga \\& Kallman 2001), and the P~Cygni lines seen from Cir~X-1 match the lines predicted by Raymond (1993) reasonably well. We also note that Iaria et~al. (2001a, 2001b) have recently found evidence for a large column density of ionized gas along the line of sight based on observations of ionized iron~K edges; the same gas might well produce both the X-ray P~Cygni lines and the ionized iron~K edges. P~Cygni lines are a common property of the ultraviolet spectra of cataclysmic variables (CVs) possessing an accretion disk. In these systems, the line profiles are associated with high-velocity polar outflows (e.g., C\\'ordova \\& Howarth 1987), in contrast to the equatorial geometry proposed for the accretion-disk wind in Cir~X-1. The proposed equatorial geometry was motivated by evidence that the Cir~X-1 system is viewed in a relatively edge-on manner. Although eclipses are not observed, the spectral variability caused by the observed X-ray absorption is best explained by a model with a relatively edge-on accretion disk (Brandt et~al. 1996; Shirey et~al. 1999b). In this model, X-rays created near the neutron star reach the observer via two light paths: (1) a direct, but often absorbed, light path that can intersect the outer bulge of the accretion disk, and (2) an indirect, electron-scattered light path that avoids the absorption associated with the direct light path. This model can naturally explain the large observed changes in the column density of the absorbing gas without corresponding changes in its apparent covering fraction (see Brandt et~al. 1996 for details). We note that radio imaging of SS433, which shares several similarities with Cir~X-1, provides strong evidence for an equatorial wind-like outflow from this system (Blundell et~al. 2001). In this second paper, we present results from a variability study of the X-ray P~Cygni lines from Cir~X-1. As in Paper~I the observations were performed with the High-Energy Transmission Grating Spectrometer (HETGS) onboard \\chandra at ``zero phase.'' Zero phase in Cir~X-1 is thought to be associated with the periastron passage of the neutron star, and near this phase Cir~X-1 is highly variable in X-ray intensity. We present analyses of two observations, which include a re-assessment of the data presented in Paper~I representing a high-flux state of Cir~X-1 and a second observation during which Cir~X-1 was in a low-flux state with strong variability. In Schulz \\& Brandt (2001; hereafter SB2001) we already presented a broad-band continuum fit to some of the data as well as showed temporal changes of the Si~XIV line. The analysis here focuses on a comparison of the bright lines in the two zero-phase flux states and their variability with intensity and time. Throughout this paper, we adopt a distance to Cir~X-1 of 6~kpc (Stewart et~al. 1993; Case \\& Bhattacharya 1998) and an interstellar column density of $2\\times10^{22}$ cm$^{-2}$ (e.g., Predehl \\& Schmitt 1995). \\centerline{\\epsfxsize=8.5cm\\epsfbox{figure1.eps}} \\figcaption{ The top panel shows the 1st order light curve of observation~II in 60~s bins. The bottom panel shows the \\xte ASM light curve including 10 days before and after observation~II was performed with \\chandra (the \\chandra observation window is marked by dotted lines). A similar plot for observation~I can be found in Paper~I.\\label{lightcurve}} ", + "conclusions": "\\subsection{Refined constraints on the accretion-disk wind model} In Paper~I we argued that the P~Cygni profiles are most likely produced by a high-velocity outflow from an accretion disk viewed in a relatively edge-on manner. The central X-ray source illuminates the disk and produces a wind driven by both thermal and radiation pressure (e.g., Begelman et~al. 1983). The intermediate-temperature ($\\sim 5\\times 10^6$~K) part of this wind produces the H-like and He-like X-ray lines (e.g., Raymond 1993; Ko \\& Kallman 1994) with P~Cygni profiles. Our current data remain generally consistent with this scenario, although we recognize that the strong wind variability seen from Cir~X-1 is not usually accounted for in the predominantly time-independent theoretical studies; we use these studies as general guides to interpretation rather than for strict comparison purposes. Our new analyses, combined with other recent work on accretion-disk winds and photoionized plasmas, have allowed us to extend some of the findings in Paper~I. For example, one of the potential problems with the disk-wind model noted in \\S3 of Paper~I can now be ameliorated. There we noted that likely values for the wind's launching radius ($r_{\\rm launch}\\approx 10^5$~km), launching density ($n_{\\rm launch}\\approxgt 10^{15}$~cm$^{-3}$), and ionization parameter ($\\xi\\lesssim 1000$~erg~cm~s$^{-1}$) led to a wind that was optically thick to electron scattering; most line photons attempting to traverse the wind would then be Compton-scattered out of the line. However, our analysis in \\S3.4, which uses the calculations of Kallman \\& Bautista (2001) rather than those of Kallman \\& McCray (1982), now suggests that a significantly higher ($\\xi\\approx 20,000$~erg~cm~s$^{-1}$) ionization parameter is likely for much of the gas creating the P~Cygni lines. The required column density along the line of sight is thus correspondingly reduced. Hydrogen column densities of $\\approx 10^{23}$~cm$^{-2}$ can now be accomodated with only moderate shielding of the wind from the full X-ray continuum or clumping of the wind. The density measurement in the line-emission region of $\\approx~2\\times 10^{13}$~cm$^{-3}$ using the Mg~XI triplet (see \\S3.3) provides a further independent constraint on the disk-wind model. Following the arguments in \\S3 of Paper~I with the new ionization parameter of $\\xi\\approx 20,000$~erg~cm~s$^{-1}$, the expected launching density is $n_{\\rm launch}\\approxgt 5\\times 10^{13}$~cm$^{-3}$. The measured density in the line-emission region is somewhat, but not greatly, lower than the launching density. Note that since the launching density depends upon the assumed launching radius, the general consistency found above supports our adopted $r_{\\rm launch}\\approx 10^{5}$~km (chosen to make the observed terminal velocity of the wind comparable to the escape velocity at $r_{\\rm launch}$; see \\S3 of Paper~I). Gratings observations of other neutron star LMXBs thus far have not generally revealed P~Cygni lines as prominent as those observed from Cir~X-1 (e.g., Cottam et~al. 2001; Marshall et~al. 2001; Paerels et~al. 2001; Schulz et~al. 2001; Sidoli et~al. 2001). It is difficult to account for this difference quantitatively, but it may be due to a combination of mass accretion rate (relative to the Eddington rate) and inclination. The high, and perhaps super-Eddington, mass accretion rate of Cir~X-1 allows an unusually powerful outflow to be driven from the accretion disk, and a relatively edge-on geometry is optimal for the viewing of an equatorial outflow. \\subsection{Line and continuum variability} Given the results from \\asca and \\xte (Brandt et~al. 1996; Shirey et~al. 1999b), the large difference in overall X-ray flux between observation~I and observation~II is likely to be due mostly to a change in the amount of absorbing material along the line of sight. As discussed by several authors (see \\S1), this material is probably associated with a thickening of the accretion disk induced by the strong mass transfer occurring at zero phase. The change in Fe~K edge depth observed by \\chandra is consistent with a large change in the amount of absorbing material, although it is difficult to constrain the exact column density during observation~II. The fact that we observe the same basic P~Cygni lines during observation~I and observation~II supports the idea that the underlying ionizing continuum source has not changed dramatically in strength. Furthermore, the fact that the line equivalent widths are significantly larger during observation~II than observation~I (see Figure~4) is generally consistent with the idea that blockage of direct continuum emission is occurring during observation~II. However, it is worth noting that the emission-line fluxes do drop significantly between observation~I and observation~II (see \\S3.3). This can be understood without invoking continuum changes if the absorbing material screens ionizing photons from reaching a significant fraction of the line-emitting gas (this screening could be stronger at longer wavelengths, explaining the trend with wavelength noted in \\S3.3). However, while it seems clear that changes in absorption play a critical role in causing much of the variability of Cir~X-1, some changes in underlying continuum strength and shape are probable as well. These are particularly likely during observation~I since it was made while Cir~X-1 was undergoing an X-ray flare above its ``base'' level (see Figure~1 of Paper~I). Our data are not ideal for probing the nature of the continuum variability in detail, but Shirey et~al. (1999a) have demonstrated that Cir~X-1 shows spectral variations that classify it as a ``Z'' type LMXB. The \\chandra data strongly suggest that much of the short-term P~Cygni line variability is caused by spectral changes in the local continuum which lead to changes in the ionization level of the wind (see \\S4), although it is difficult to rule out entirely changes in the wind geometry as well. The observed short-term changes of the lines can be dramatic, causing them to vary from being almost completely in absorption to almost completely in emission. Modeling this variability in detail will be a challenge. \\subsection{Comparisons with the spectral features seen from superluminal jet sources} The Fe~K line region in the low-flux state spectrum of Cir~X-1 shows strong, blueshifted resonance absorption features from Fe~XXV and Fe~XXVI. Significant line emission during the low-flux state is only seen from neutral or nearly neutral Fe. Such spectral behavior is reminiscent of that seen from Galactic superluminal jet sources containing a black hole like GRS~1915+105 (Kotani et~al. 2000; Lee et~al. 2001) and GRO~J1655--40 (Ueda et~al. 1998; Yamaoka et~al. 2001). Similar Fe absorption has also now been observed in GX~13+1 (Ueda et~al. 2001), a neutron star LMXB that has been classified as an atypical atoll source (Hasinger \\& van der Klis 1989), and MXB~1659--298 (Sidoli et~al. 2001), also probably an atoll source; these results show that these spectral features can appear independent of the nature of the compact object and the existence of jets. In most of these cases it has been argued that the absorption lines come from a highly ionized plasma layer extending to large heights above the equatorial plane (Yamaoka et~al. 2001). Specifically, a geometrically thick accretion flow at $\\sim 10^4$ Schwarzschild radii is assumed, which in the case of GRO~J1655--40 amounts to $\\sim 10^5$ km. This is of the same magnitude as our launching radius for the wind in Cir~X-1, and it is worth considering the possibility that the material creating the Fe X-ray absorption lines in superluminal sources might generally be outflowing rather than inflowing. An equatorial wind scenario appears consistent with the gratings observations of GRS~1915+105 (Lee et~al. 2001); if the outflow velocity for GRS~1915+105 is assumed to be the same as that for Cir~X-1, the outflow would need to be less inclined relative to the line of sight to be consistent with the 770~km~s$^{-1}$ upper limit on the blueshift. The strong Fe~XXVI absorption feature from GRS~1915+105 shows variations on similar time scales to those observed from Cir~X-1. Superluminal jet sources containing black holes have not shown as many spectral features as we observe from Cir~X-1; typically only features from highly ionized iron are seen. This could plausibly be explained if their winds are more highly ionized due to their larger luminosities and different continuum shapes, so that iron is the only abundant element not fully stripped of its electrons." + }, + "0112/astro-ph0112160_arXiv.txt": { + "abstract": "We observed Cyg X-1 with {\\em RXTE} contiguously over its 5.6-day binary orbit. The source was found to be in the hard state throughout the observation. Many intensity dips were detected in the X-ray light curves. We found that the dips fell into two distinct categories based on their spectral properties. One type exhibits strong energy-dependent attenuation of X-ray emission at low energies during a dip, which is characteristic of photoelectric absorption, but the other type shows nearly energy-independent attenuation. While the first type of dips are likely caused by density enhancement in an inhomogeneous wind of the companion star, as previous studies have indicated, the second type might be due to partial obscuration of an extended X-ray emitting region by optically thick ``clumps'' in the accretion flow. It is also possible that the latter are caused by a momentary decrease in the X-ray luminosity of the source, due, for instance, to a decrease in the mass accretion rate, or by Thomson scattering in highly ionized ``clumps''. We discuss the implications of these scenarios. ", + "introduction": "Cyg X-1 was the first astronomical system discovered to show strong evidence for a stellar mass black hole (Bolton 1972; Webster $\\&$ Murdin 1972). It is a binary system with an orbital period of 5.6 days. The results from radial velocity measurements imply that the mass of the compact object exceeds the upper limit on the mass of a neutron star, so the compact object is inferred to be a black hole. A more recent study, based on spectrum synthesis, has derived a mass of about 10 M$_{\\odot}$ for the black hole (Herrero et al. 1995). The companion star has been identified as an O9.7 Iab supergiant (Walborn 1973; Gies $\\&$ Bolton 1986), with a mass of about 20 M$_{\\odot}$ (Herrero et al. 1995). Cyg X-1 is, therefore, intrinsically different from most known black hole candidates which contain only a low-mass companion star. This difference might be related to the fact that Cyg X-1 is a persistent X-ray source while those with a low-mass companion are exclusively transient sources. The X-ray emission from Cyg X-1 is likely powered by accretion of material from the companion star by the black hole. In this case, the accretion flows are thought to follow a pattern that is intermediate between that of Roche-lobe overflow, as in low-mass systems, and that of wind accretion, which is common for high-mass systems (Gies $\\&$ Bolton 1986). Such an accretion process is sometimes referred to as ``focused wind accretion'', which can occur when the companion star is close to filling its Roche lobe (Friend \\& Castor 1982). The X-ray observations of Cyg X-1 have revealed that it has two distinct spectral states, hard and soft (see reviews by Oda 1977, Liang $\\&$ Nolan 1984, Tanaka \\& Lewin 1995, Cui 1998, and Liang 1998). The source is usually found in the hard state, when its X-ray spectrum is relatively flat (with a typical power-law photon index of 1.5), but, occasionally, it makes a transition to the soft state, when the spectrum steepen significantly (to a typical photon index of 2.5). In the hard state, the X-ray intensity of Cyg X-1 is strongly modulated by the binary motion (Wen et al. 1999; Brocksopp et al. 1999; Priedhorsky, Brandt, $\\&$ Lund 1995; Holt et al. 1979), probably caused by varying amount of absorbing material along the line of sight through the stellar wind of the companion star. As expected, the intensity is found to reach a minimum at the times of superior conjunction of the black hole (when the companion star is in front of the black hole with respect to the line of sight), which is usually defined as orbital phase $0$. It is, however, a very broad minimum, spanning about 26$\\%$ of the orbit. In contrast, no X-ray orbital modulation is measurable in the soft state (Wen et al. 1999), which might be due to a significant change in the physical conditions of the wind (e.g., much higher ionization level) and/or in the geometry of the X-ray emitting region. Besides the global orbital modulation of the X-ray emission, X-ray intensity dips are often seen in the light curves of Cyg X-1 (e.g., Pravdo et al. 1980; Remillard \\& Canizares 1984; Kitamoto et al. 1984; Ba{\\l}uci{\\'n}ska \\& Hasinger 1991; Ebisawa et al. 1996). The dips vary in duration from minutes to hours. The spectrum of the source hardens during a dip, implying that the dips probably originate in the photoelectric absorption of X-rays by the ambient medium. This is further supported by the detection of Fe K absorption edge in the dip spectrum in some cases (Kitamoto et al. 1984). The absorbing column density can increase by more than one order of magnitude during a dip (Kitamoto et al. 1984; Ba{\\l}uci{\\'n}ska-Church et al. 1997). The distribution of dips against orbital phases shows a prominent peak around phase $0$ (Ba{\\l}uci{\\'n}ska-Church et al. 2000), implying an origin of the dips in the stellar wind from the companion star. In this paper, we present results from a long observation of Cyg X-1 with the {\\em Rossi X-Ray Timing Explorer} (RXTE) over one complete 5.6-day orbital cycle. Many intensity dips were detected. The large collecting area of the Proportional Counter Array (PCA) aboard {\\em RXTE} made it possible to conduct a more detailed investigation of spectral evolution of the source during some of the strong dips, as well as to carry out orbital-phase-resolved spectroscopy to quantify the variation in the column density over an orbital cycle. ", + "conclusions": "As mentioned in the introduction, the phenomenon of X-ray intensity dips in Cyg X-1 is well known and well studied. Over the years, a rich database has been built up for this famous black hole candidate. The database is essential for a systematic study of dips, whose occurrence is somewhat unpredictable and is thus easy to miss in a brief pointed observation. The distribution of the dips over a binary orbital cycle has been reliably established recently (Ba{\\l}uci{\\'n}ska-Church et al. 2000): most of the dips occur near superior conjunction of the X-ray source, but dipping activity has also been seen at orbital phases far from superior conjunction. Our results are consistent with that. We went a step further in this study. We carried out a long observation of Cyg X-1 over one entire orbital period with the large-area detectors aboard {\\em RXTE}. The data allowed us to conduct more detailed spectral analyses of the dipping phenomenon as well as phase-resolved spectroscopy to quantify the variation in the column density over an orbit. The main result of this investigation is the recognition of the existence of two types of dips. The spectrum of type A dips shows an energy-dependent reduction at low energies that is characteristic of photoelectric absorption. Detailed spectral modeling shows that the only difference between the spectrum of type A dips and the average non-dip spectrum seems to be the presence of additional column density during a dip. We stress, however, that our data is not of sufficient quality to warrant more complicated modeling, as has been done previously (e.g., Ebisawa et al. 1996). The type A dips seem to occur preferentially around the times of superior conjunction of the black hole. The spectrum of type B dips, on the other hand, shows almost an energy-independent reduction at low energies. Therefore, these dips cannot possibly be due to photoelectric absorption. Moreover, the type B dips appear to distribute randomly over the binary orbit. Type A dips are much more common than type B dips. They are probably produced by density enhancement in an inhomogeneous wind from the companion star, since the column density changes only moderately during such a dip. They occur more frequently around superior conjunction probably because the line-of-sight follows a longer path through the denser part of the wind (however, see Blondin \\& Woo 1995). The inhomogeneities in the wind can be caused by a variety of physical processes (see Ba{\\l}uci{\\'n}ska-Church et al. 2000 for a few examples). It remains to be seen whether the observed global orbital modulation of X-ray emission is entirely due to the presence of such dips. Type B dips might be caused by partial covering of an extended X-ray emission region by an opaque ``screen''. In the context of Comptonization models (e.g., review by Liang 1998 and references therein), the presence of an extended emitting region may not be unreasonable for Cyg X-1. The question is what could physically serve as an opaque ``screen'' to produce the observed type B dips. The answer may lie in the fact that Cyg X-1 is a high-mass X-ray binary (HMXB), in which the mass accretion process is probably mediated by stellar wind from the massive companion star. Numerical simulations of mass accretion in HMXBs show that a tidal stream develops when the companion star is close to filling its Roche lobe and that the stream trails behind the compact object due to the Coriolis force (Blondin et al. 1991). The density in the stream is shown to be as high as 20--30 times the ambient wind density. Since the obscuration of X-rays occurs close to the accretion disk in this case, it is not expected to be very sensitive to the binary motion. Now, the question is whether this scenario can be applied to Cyg X-1, whose orbit is only moderately inclined, with a probable inclination angle of 30\\arcdeg -- 40\\arcdeg\\ (Bolton 1975; Guinan et al. 1979; Daniel 1981). A full three-dimensional simulation would be required to show the scale height of the tidal streams. We note that the inclination angle remains to be poorly determined for Cyg X-1, which is one of the largest contributors to the uncertainty in determining the mass of the compact object. For instance, the optical polarization measurements imply a wide range of 25\\arcdeg -- 70\\arcdeg (e.g., Long et al. 1980). Type B dips might also be due to a sudden decrease in the mass accretion rate, causing a decrease in the X-ray luminosity of the source itself. This could, in principle, occur, since, after all, the intrinsic variability of the source is generally thought to be due to fluctuation in the mass accretion rate. However, for Cyg X-1 (and black hole candidates in general), a change in X-ray flux is generally accompanied by a change in spectral hardness. In fact, it has been shown that in the hard state the flux and spectral hardness are, to some degree, anti-correlated (e.g., Wen et al. 2001). However, the correlation is quite weak, so our results cannot rule out this possibility. It is also possible that type B dips are produced by (nearly) pure Thomson scattering of X-rays in highly ionized dense clumps in the wind. Given that the cross section for photoelectric absorption is, on average, more than two orders of magnitude larger than the Thomson cross section, the density of a clump that produces a type B dip is required to be more than two orders of magnitude larger than that of a clump that produces a type A dip of comparable depth (assuming that the clumps are comparable in size). To abtain an ionization parameter that is, for instance, a factor of 100 larger for type B clumps than for type A clumps, we would then require that the former are, on average, more than 100 times closer to the X-ray source. If so, this scenario could also explain why the type B dips are much less dependent of orbital phases. Finally, we made an attempt to quantify how the column density varies with orbital motion by carrying out orbital-phase-resolved spectroscopy. The results seem to indicate that the column density is higher at superior conjunction, although error bars are too large for us to be definitive about it. The issue can be resolved by a similar observation with detectors of much improved low-energy response and spectral resolution, such as those aboard {\\em Chandra} and {\\em XMM}. Much improved spectral resolution of these detectors will also allow the detection of absorption edges, which provides means of determining column density in a manner that is independent of continuum modeling (Schulz et al. 2001). Furthermore, the highly eccentric orbits of these satellites allow long un-interrupted observations, which is very much important for studying the distribution of dips over a binary orbit." + }, + "0112/astro-ph0112430_arXiv.txt": { + "abstract": "The Large Zenith Telescope Survey whose construction is almost completed (first light expected in spring 2002) near Vancouver (Canada) is designed to observed a total strip of $\\sim17\\arcmin\\times 120^\\circ$ in 40 narrow-band filters spanning 4000-10000 \\AA. It will gather the spectrophotometric energy distributions of ca. $10^6$ galaxies to redshifts $z\\sim1$, with redshift accuracy $\\sigma_z=0.01$ at s/n=10, $\\sigma_z=0.04$ at s/n=3, ca. $10^5$ stars, and a large sample of QSOs, variable stars, and transient objects of the solar system. The survey is optimized for studying of the evolution of both the luminosity function and the clustering of galaxies to a redshift $z\\sim1$. It will also provide a complete and homogeneous sample of stars at various galactic latitudes useful for studying galactic structure, and it will be a good instrument for the monitoring of variable objects. ", + "introduction": " ", + "conclusions": "" + }, + "0112/astro-ph0112095_arXiv.txt": { + "abstract": "We present discovery spectra of a sample of eight lensed galaxies at high redshift, $3.7<$z$<5.2$, selected by their red colors in the fields of four massive clusters: A1689, A2219, A2390, and AC114. Metal absorption lines are detected and observed to be blueshifted by 300-800 km s$^{-1}$ with respect to the centroid of Ly-$\\alpha$ emission. A correlation is found between this blueshift and the equivalent width of the metal lines, which we interpret as a broadening of saturated absorption lines caused by a dispersion in the outflow velocity of interstellar gas. Local starburst galaxies show similar behavior, associated with obvious gas outflows. We also find a trend of increasing equivalent width of Ly-$\\alpha$ emission with redshift, which may be a genuine evolutionary effect towards younger stellar populations at high redshift with less developed stellar continua. No obvious emission is detected below the Lyman limit in any of our spectra, nor in deep $U$ or $B$-band images. The UV continua are reproduced well by early B-stars, although some dust absorption would allow a fit to hotter stars. If B-stars dominate, then their relatively prominent stellar absorption lines should separate in wavelength from those of the outflowing gas, requiring more detailed spectroscopy. After correcting for the lensing, we derive small physical sizes for our objects, $\\sim$ 0.5-5 kpc h$^{-1}$ for a flat cosmology with $\\Omega_m=0.3, \\Omega_{\\Lambda}=0.7$. The lensed images are only marginally resolved in good seeing despite their close proximity to the critical curve, where large arcs are visible and hence high magnifications of up to $\\sim20\\times$ are inferred. Two objects show a clear spatial extension of the Ly-$\\alpha$ emission relative to the continuum starlight, indicating a ``breakout'' of the gas. The sizes of our galaxies together with their large gas motion suggests that outflows of gas are common at high redshift and associated with galaxy formation. ", + "introduction": "The photometric selection of high redshift galaxies by their Lyman-limit break has proven to be a successful means of obtaining galaxy redshifts in the range $2.5<$ z $<3.3$ \\citep{Steidel:92, Steidel:96a, Steidel:96b, Lowenthal:97}. At somewhat higher redshift, photometric selection is less reliable because the forest depression of the continuum between Ly-$\\alpha$ and the Lyman limit can be confused with evolved or dusty galaxies at lower redshift. Moreover, these Lyman-depressed galaxies at z$\\sim$4 are significantly fainter than the Ly-break galaxies, making spectroscopic work challenging. Despite this, a sample of 48 $B$-band dropout selected galaxies has been spectroscopically confirmed by \\citet{Steidel:99}. These are observed with relatively low spectral resolution, sufficient for determining redshifts in the range $3.8<$ z $<4.5$, and occur with a frequency consistent with little evolution relative to the lower redshift Ly-break galaxies \\citep{Steidel:99}. Only a handful of higher redshift galaxies are known, the best examples being serendipitous discoveries boosted by lensing \\citep{Warren:96, Franx:97, Trager:97, Frye:98}. Others are identified primarily on the basis of an asymmetric emission line taken to be Ly-$\\alpha$ \\citep{Dey:98}, or by virtue only of a relatively large lower limit for the equivalent width of single emission lines \\citep{Weymann:98, Hu:98, Hu:99, Stern:00b, Rhoads:00, Manning:00}. In principle, high quality photometry alone is sufficient for z$>$4.5 candidates because the strength of the Lyman-series forest depression \\citep{Madau:95} can be established from bright QSO spectra \\citep{Schneider:89,Storrie:96}, yielding precise predictions about the galaxy colors at those redshifts. This continuum suppression feature has led to convincing claims of z$>$5 for the very reddest objects in the Hubble Deep Field \\citep{Lanzetta:96,Spinrad:98}, and more recently, to the successful selection of high redshift QSOs by \\citet{Stern:00} and the Sloan Digital Sky Survey collaboration \\citep[and references therein]{Fan:00,Anderson:01}. Here we make use of this established high redshift Lyman-series depression to color select galaxies for spectroscopy near the critical curves of some of the most massive lensing clusters. This approach provides a viable way of obtaining useful spectroscopy of high-z galaxies, under favorable circumstances, where lensing is compounded by a massive cluster galaxy lying close to the critical curve of a cluster \\citep{Franx:97,Frye:98}. The galaxies presented here form part of a systematic redshift survey of red-selected galaxies behind massive lensing clusters. Most galaxies selected in this way are early type lying at modest redshift behind the cluster. We present the full redshift survey, and measurements of cluster masses by magnification, in a different paper \\citep{Frye:01}. Here we show only the most distant galaxies discovered in our survey as a by-product of the red selection, and whose optical colors place them in the range 3.5$<$z$<$6.5 \\citep{Benitez:00}. We analyze the spectral information from 8 lensed galaxies detected with redshifts of z~$=3.77, 4.04, 4.07, 4.25, 4.46, 4.67, 4.87,$ and $5.12$. These objects are relatively bright, allowing for detailed spectroscopic work, and are suitable for follow-up at other wavelengths. This paper is organized as follows. In \\S2 we describe the target selection and in \\S3 we outline the photometric and spectroscopic observations and data analysis techniques. This is followed by a presentation of the photometric and spectroscopic results in \\S4, an examination of the agreement with photometric redshift results in \\S5, and an analysis of gas motion and the equivalent width of the Ly-$\\alpha$ emission line in \\S6. We close in \\S7 with a discussion of our main results. ", + "conclusions": "The claimed absence of Ly-$\\alpha$ emission in deep narrowband searches \\citep{Thompson:95} was readily explained by dust absorption along the hugely increased path length of resonantly scattered Ly-$\\alpha$ emission \\citep{Charlot:93}, for even relatively modest columns of neutral hydrogen. However, the recent discovery of copious numbers of high redshift galaxies at 2.5$<$z$<$3.5, half of which show Ly-$\\alpha$ in emission \\citep{Steidel:99,Rhoads:00}, revises this conclusion. The explanation for the escape of Ly-$\\alpha$ emission is plausibly explained as a dynamical effect from outflowing gas, based on detailed UV spectroscopy of nearby starburst galaxies \\citep{Legrand:97}. This outflow model, although not yet quantified, can account for the peculiar observations reported here and for other high quality spectra of high redshift galaxies. The Ly-$\\alpha$ photons generated by hot stars in the center of HII regions can escape if back scattered from the inner ``wall'' of the surrounding expanding gas, thereby avoiding the fate of the average Ly-$\\alpha$ photon which scatters and then is absorbed by dust within the surrounding gas. These redder, back scattered photons take on the outflow velocity. If great enough, this velocity will exceed in wavelength the redward wing of Ly-$\\alpha$ absorption of the foreground gas through which these photons must pass in order to reach the observer. In contrast, the bluer photons will be preferentially absorbed, leading to an asymmetric emission line. The internal gas motion implied by the mismatch of interstellar absorption lines with emission lines of Ly-$\\alpha$ and non-resonance lines of OI detected in the IR has established, by analogy with local starburst galaxies, that gas outflows are ubiquitous at early times. If this gas escapes the potential of these early galaxies, then the consequences for galaxy formation are predicted to be rather important. The gas belonging to neighboring density perturbations which have not yet virialized will be only tenuously held and be expected to be stripped away by winds from nearby collapsed starforming galaxies \\citep{Scannapieco:01a}. Since galaxy formation is predicted to be much more spatially correlated at early times by biasing \\citep{Cole:89}, the effect locally is to suppress the formation of small galaxies \\citep{Scannapieco:01b}. At later times, this metal enriched gas is expected to become bound to larger density perturbations which take longer to collapse and will be heated and further enriched, modifying the cooling times of massive galaxies and clusters \\citep{Scannapieco:01b}. The ubiquitous presence of metals in the IGM detected out to the highest redshifts that the forest has been probed \\citep{Frye:93,Lu:96a,Lu:96b}, and the high level of enrichment of cluster gas, may be most easily explained by a general pollution by early galaxies \\citep{Scannapieco:01a,Madau:01,Pettini:01} though other mechanisms have been proposed to contribute, including tidal stripping of processed gas \\citep{Gnedin:97}, and photoevaporation of gas in sufficiently small galaxies by UV background \\citep{Barkana:00b}. Lensing helps to spatially resolve galaxies by stretching their images so that one can obtain intrinsic size measurements. We have made this calculation for several of our highly magnified cases, which after the correction for the likely level of magnification show that the objects are in fact very small, $\\sim$0.5-5~kpc h$^{-1}$. This is in reasonable agreement with a sample of very distant red objects identified in the HDF fields \\citep{Spinrad:98, Weymann:98}. It also agrees with theoretical predictions \\citep{Barkana:00a} for the early evolution expected in the hierarchical models for structure formation, in which the first galaxies to form are small, with increasingly more massive objects collapsing and merging over time. We conclude that the lensed galaxies presented here provide high quality information on the earliest known galaxies. Lensing has fortuitously afforded us detailed spectral and spatial details of magnified but presumably otherwise typical examples of galaxies at z$>$4. The rarity of luminous high redshift galaxies and the requirement of a high magnification for useful spectroscopic follow-up means that such work will be slow but well rewarded with precious data on galaxies at otherwise inaccessibly early times." + }, + "0112/astro-ph0112048_arXiv.txt": { + "abstract": "In cataclysmic variables (CVs), accretion onto white dwarfs produces high temperature, high density plasmas. They cool down from kT$\\sim$10 keV via bremsstrahlung continuum and K and L shell line emissions. The small volume around white dwarfs means that the plasma densities are much higher than in, e.g., stellar coronae, probably beyond the range well-described by existing models. I will describe potential diagnostics of the temperatures, the densities, and the optical depths of X-ray emitting plasmas in CVs, and present the recent Chandra grating spectra of the magnetic CV V1223 Sgr as an example. ", + "introduction": " ", + "conclusions": "" + }, + "0112/astro-ph0112562_arXiv.txt": { + "abstract": "{ We present a power spectrum analysis of the long \\asca observation of \\ark in June/July 2001. The observed power spectrum covers a frequency range of $\\sim 3.5$ decades. We detect a high frequency break at $\\sim 2\\times 10^{-3}$ Hz. The power spectrum has an rms of $\\sim 30\\%$ and a slope of $\\sim -1$ and $\\sim -2$ below and above the break frequency. When combined with the results from a long \\xte observation (Pounds \\etal 2001), the observed power spectra of \\ark and Cyg X-1 (in the low/hard state) are almost identical, showing a similar shape and rms amplitude. However, the ratio of the high frequency breaks is very small ($\\sim 10^{3-4}$), implying that these characteristic frequencies are not indicative of the black hole mass. This result supports the idea of a small black hole mass/high accretion rate in \\ark. ", + "introduction": "\\smallskip Narrow-line Seyfert~1 (NLS1) galaxies are a peculiar group of active galactic nuclei (AGN) characterized by their distinct optical line properties (Osterbrock \\& Pogge 1985). In hard X-ray studies NLS1 galaxies comprise less than 10\\% of the Seyfert galaxies, however, from the ROSAT All-Sky Survey it became clear that about half of the AGN in soft X-ray selected samples are NLS1 galaxies (Grupe 1996, Hasinger 1997). Boller \\etal (1996) found from ROSAT observations that the soft X-ray spectra of NLS1 galaxies are systematically steeper than those of broad line Seyfert~1 galaxies. They further discovered that NLS1 galaxies frequently show rapid short time scale X-ray variability which can be interpreted as evidence for a small black hole masses in these objects. Ark~564 is the X-ray brightest NLS1 galaxy with a 2$-$10~keV flux of $\\sim 2\\times10^{-11}$ \\flux (Vaughan \\etal 1999) and shows large amplitude variations on short time scales (Leighly 1999). Therefore it is the best candidate to study its X-ray variability in order to obtain important clues about the size of the black hole mass and the accretion rate in NLS1 galaxies. Recently, \\ark was observed by \\xte once every $\\sim 4$ days from January 1999 to September 2000 covering $\\sim 20$ months of data. Its $2-10$ keV power spectral density (PSD) function showed a cut--off at a frequency which corresponds to a time scale of $\\sim 13$ days (Pounds \\etal 2001). \\ark was also observed for a period of $\\sim 35$ days in June/July 2000 by \\asca. This observation was part of a multi-wavelength AGN Watch monitoring campaign (Turner \\etal 2001). The flux variations that \\ark exhibited during this observation have already been studied by Edelson \\etal (2001) who found that the variability amplitude is almost independent of energy band and the power spectrum is harder in the hardest energy bands. There were no delays between the variations in the different bands and no signs of non-linear behaviour in the light curves. Gliozzi \\etal (2001) also found no sign of non-linear behaviour and no statistically significant indication of non-stationarity in the light curves. Furthermore, using nonlinear techniques they were able to demonstrate that the source behaves differently in the high and low flux states. In this work we present a timing analysis of the June-July \\asca observation of \\ark based on power spectral analysis techniques. In the following section we present the data and the results from the PSD analysis while in Section 3 we discuss our conclusions. ", + "conclusions": "Using the \\asca June-July 2001 observation of \\ark we have estimated its power spectrum over a frequency range of $\\sim 3.5$ decades. We have detected a high frequency break at $\\sim 2\\times 10^{-3}$ Hz. In Figure 3, together with the $0.7-5$ keV PSD, we have also plotted the best fitting ``cut-off power law\" model from Pounds \\etal (2001). The observed PSD of \\ark extends over $\\sim 6$ decades in frequency. It is clear from this Figure that the PSD starts to flatten below $\\sim 10^{-5}$ Hz, showing a second, low frequency break. MCG 6-30-15 also shows two break frequencies in its PSD (Nowak \\& Chiang 2000), however, the evidence for the low frequency break is not significant in that case (Uttley \\etal 2001). Therefore, \\ark is probably the only Seyfert galaxy for which we are certain that the observed PSD has a shape that is almost identical to the shape of the Cyg X-1 PSD in the low/hard state. In both cases, the PSD slope is flat below a characteristic low frequency. Above this frequency the PSD steepens to $\\propto \\nu^{-1}$ and then becomes $\\propto \\nu^{-2}$ above a second break frequency. Apart from the PSD shape, the rms amplitudes at frequencies higher than the low frequency cut off are also very similar. These results show that \\ark corresponds to the low/hard (instead of the high/soft) state of the Galactic black hole candidates (GBHC). However, despite this similarity, \\ark is not simply a ``larger\" version of GBHCs. The two break frequencies in Cyg X-1 are located between $0.03- 0.3$ and $1-10$ Hz (e.g. Nowak \\etal 1999, Negoro \\etal 2001). The ratio between the characteristic frequencies in \\ark and Cyg X-1 is therefore $\\sim 10^{5-6}$ and $\\sim 10^{3-4}$, if we consider the low and high frequency breaks, respectively. Although there is an uncertainty associated with these ratios (for example, we do not know whether the break frequencies in AGN are constant or change with time) this cannot probably explain the two orders of magnitude difference between the ratios. Perhaps then the break frequencies do not scale proportionally with the black hole mass. For example, if the high frequency break in Cyg X-1 corresponds to the Keplerian orbital period at $\\sim 30 R_{S}$ around the central object (Nowak \\etal 1999; $R_{S}$ is the Schwarzschild radius) and the high frequency break in \\ark corresponds to the same time scale but at a radius $\\sim 3 R_{S}$, then the ratio of the characteristic time scales will not be equal to the mass ratio but to the mass ratio $\\times (30 R_{S}/3 R_{S})^{-3/2}$. On the other hand, if the low frequency breaks scale with mass, then the central black hole mass of \\ark should be $\\sim 10^{7}$ M$_{\\odot}$, assuming that the black hole in Cyg X-1 is $10$ M$_{\\odot}$ (Herrero et al. 1995). If that is the case, taking into account the luminosity of \\ark, Pounds \\etal (2001) concluded that the source is accreting at a substantial fraction of the Eddington limit. The orbital period at $3R_{S}$ for a $10^{7}$ M$_{\\odot}$ black hole is $\\sim 5000$ sec, which is an order of magnitude larger than the time scale which corresponds to the high frequency break in the PSD of \\ark. Either the mass of the black hole in \\ark is smaller than $10^{7}$ M$_{\\odot}$ (in which case the accretion rate is larger than the Eddington limit), or due to the large accretion rate, the characteristic time scales are smaller. When the accretion rate approaches the Eddington limit advective energy transport dominates over radiative cooling and the disk becomes moderately geometrically thick, a so called ``slim disk\" (Abramowicz et al. 1988). In this case, a substantial amount of radiation can be produced from inside the last stable orbit, i.e. $3R_{S}$ (Mineshige et al. 2000). The free-fall time scale at distance $R$ is given by $t_{\\rm ff}=(R_{G}/c)(R/R_{G})^{3/2}=500$M$_{8}(R/R_{G})^{3/2}$ sec (Rees 1984), where M$_{8}$ is the mass in units of $10^{8}$M$_{\\odot}$ and $R_{G}=GM/c^{2}$ is the gravitational radius. For M$_{8}=0.1$ and $R=3R_{S}=6R_{G}$, we get $t_{\\rm ff}\\sim 700$ sec which is comparable to $1/\\nu_{\\rm bf}$ that we find for \\ark. The sound crossing time scale is another possible candidate that could correspond to $1/\\nu_{\\rm bf}$. In the thin disk case this time scale is very long even in the innermost parts of the disk. However, when the accretion rate is comparable to the Eddington rate and the disk is dominated by radiation pressure, the group velocity of sound waves (responsible for transmitting density fluctuation information) becomes very large (e.g. Krolik et al. 1991). Consequently, the sound crossing time scale at say $R=3R_{S}$ will be decreased substantially and could be of the order of $1/\\nu_{\\rm bf}$. We conclude that, the similarity of the PSD and time lags in \\ark implies that X-rays in this source are produced by a mechanism similar to the X-ray emission mechanism in Cyg X-1 in its low/hard state. On the other hand, the comparison between the characteristic time scales in the two systems suggests that at least one physical parameter is different in them, probably the accretion rate, with \\ark having a substantially higher accretion rate than Cyg X-1 in its low/hard state. \\vskip 0.4cm" + }, + "0112/astro-ph0112081_arXiv.txt": { + "abstract": "We present a new particle based code with a multi-phase description of the ISM implemented in order to follow the chemo-dynamical evolution of galaxies. The multi-phase ISM consists of clouds (sticky particles) and diffuse gas (SPH): Exchange of matter, energy and momentum is achieved by drag (due to ram pressure) and condensation or evaporation. Based on time scales we show that in Milky-Way-like galaxies the drag force is for molecular clouds only important, if their relative velocities exceed 100 km/s. For the mass exchange we find that clouds evaporate only if the temperature of the ambient gas is higher than one million Kelvin. At lower temperatures condensation takes place at time scales of the order of 1--10~Gyr. ", + "introduction": " ", + "conclusions": "" + }, + "0112/astro-ph0112424_arXiv.txt": { + "abstract": "To confront the predictions of the most recent line-driven disk wind models with observational evidence, we have obtained \\hst\\ STIS (1180--1700\\thinspace\\AA) echelle spectra of the nova-like variables IX~Vel and V3885~Sgr at three epochs. The targets were observed in timetag mode for $\\sim$2000\\thinspace sec on each occasion, allowing us to study the spectral time evolution on timescales down to $\\sim$10\\thinspace sec. The mean UV spectra are characterised by the wind signature of broad blueshifted absorption in \\lya, N\\Vl1240, Si\\IVl1398, C\\Vl1549 and He\\IIl1640. There is very little redshifted emission other than in C\\IV. Narrow blueshifted absorption dips, superposed on the broad absorption at around $-900$\\kms, accompany periods of well-developed wind activity. The continuum level and mean line profiles vary markedly from observation to observation -- with the wind signatures almost disappearing in one epoch of observation of IX~Vel. The strong positive correlation between UV brightness and wind activity predicted by line-driven disk wind models is disobeyed by both binaries. The wind signatures in IX~Vel's UV spectrum are revealed to be remarkably steady on timescales ranging from $\\sim$10 to $\\sim$1000\\thinspace sec. More variability is seen in V3885~Sgr, the binary with the lower opacity outflow. But there is only one epoch in which the line profile changes significantly in $\\sim$100\\thinspace sec or less. Narrow absorption dips, when present, show only smooth, small changes in velocity. We surmise these may trace the white dwarf's orbital motion. The near-absence of line profile variability on the shorter 10- to 100-sec timescales, and the lack of correlation between wind activity and luminosity, could both arise if a non-radiative factor such as the magnetic field geometry controls the mass loss rate in these binaries. ", + "introduction": "Accretion disks are present in many astrophysical systems, with firm evidence for their existence in environments as diverse as those of young stars \\cite{98hartmann}, close binary systems and AGN (Frank, King \\& Raine 1985). Objects known to include hot optically-thick disks typically provide clear spectroscopic evidence of mass loss in the form of broad blueshifted UV absorption features. Currently there are two main driving mechanisms envisaged for disk winds. The more widely studied of these, regarded as the only option for low-mass YSOs, relies on MHD processes (e.g. Ouyed \\& Pudritz 1997, Hirose et al. 1997, Konigl \\& Pudritz 2000). In high luminosity systems an alternative or additional method for driving an outflow is radiation pressure mediated by line opacity (hence the term `line-driven wind'). Until recently the multi-dimensional modelling required to deal with the geometry of a line-driven disk wind was beyond reach. However, this problem has now been approached numerically, with greater success, in the models of Pereyra, Kallman \\& Blondin (1997) and Proga, Stone \\& Drew (1998, hereafter PSD). As a result of this work, the radiation-driven wind model already provides some clear-cut predictions that can be put to the test against observations. Pre-eminent among these is the expectation that the wind mass loss rate will show a strong dependence on system luminosity (see Drew \\& Proga 2000). PSD also predict unstable, clumpy outflow in the regime where the accretion disk, rather than the central object, contributes most of the radiant luminosity. Indeed, PSD's simulations suggested that the dominant mass loss stream flowing away from the inner disk will include denser clumps of gas that fail to reach escape velocity and collapse back onto the disk at larger radii. A better treatment of the line-driving (Proga, Stone \\& Drew 1999) preserves this impression of a highly structured, clumpy outflow. When observed spectroscopically with a time resolution of less than the typical flow timescale this phenomenon could make itself apparent as narrow absorption components, superposed on the broad blueshifted absorption profile, sweeping both away from and in towards line centre. With the aim of looking for the spectral signatures of disk-wind variability on suitably short timescales we have obtained {\\it Hubble Space Telescope} (\\hst) observations of two high-state cataclysmic variables, using the Space Telescope Imaging Spectrograph (STIS). Specifically, our targets are the nova-like variables: IX Vel and V3885 Sgr. These binaries provide examples of disk winds where the line-driving mechanism may dominate over MHD processes in powering mass loss. Both IX~Vel and V3885~Sgr are made up of a low mass star overflowing its Roche lobe and transferring matter to a non- or weakly-magnetic white dwarf (WD) via an accretion disk (i.e. they are UX~UMa type cataclysmic variables). Both are known to persist in a high brightness state, powered by high mass accretion rates ($\\la10^{-8}$\\msunyr). This puts these objects just within the domain where disk luminosities are $\\sim0.001L_{\\rm{E}}$ (where $L_{\\rm{E}}$ is the Eddington luminosity) and spectral line opacity may be sufficient to power significant mass loss \\cite{2000drew}. Essentially all of the light from these systems at optical and ultraviolet wavelengths is due to accretion. A further factor influencing target choice is that IX~Vel and V3885~Sgr are the brightest of the known UX~UMa binaries. With the possibility of high-quality observation of the wind-shaped UV line profiles, these binaries can provide a particularly good test bed for line-driven disk wind models. On timescales of weeks and months, IX~Vel and V3885~Sgr vary by several tenths of a magnitude. In order to allow this secular variation to provide a sampling of a modest range of luminosity states , both targets were observed three times each at intervals of at least a month. Both targets are non-eclipsing systems. The broad blueshifted absorption features seen in the UV, that are indeed a common feature of the lower inclination UX UMa systems, are not usually very deep, making them ideal for studying the fine structure of wind-formed lines. To obtain the high time and wavelength resolution required for this programme, STIS was employed in its time-tag observing mode, using an echelle to disperse the light. This allows the observer to select the desired wavelength or time resolution at the point of data extraction and calibration. To ensure a minimum S/N ratio there is necessarily a trade-off such that higher time resolution will result in a reduced useable spectral resolution and vice-versa. This is not the first use of the \\hst\\ spectrographs for high time resolution studies of cataclysmic variable disk winds: Prinja et al. (2000a), (2000b) report results of Goddard High Resolution Spectrometer (GHRS) grating observations of respectively BZ~Cam and V603~Aql. These data gave a time resolution of $\\sim$80 and $\\sim$30 seconds, timescales that are also comparable with the likely outflow timescales in these binaries (of order a few tens up to $\\sim$100 seconds). Significant absorption line profile variability on timescales down to $\\sim60$~sec was reported. The key differences between these published studies and the present study lie in data quality -- our brighter targets allow the use of the higher resolution echelle -- and in the prior knowledge of the targets' UV spectral characteristics. The organisation of this paper is as follows. First, we present a brief description of the STIS/time-tag observations (section \\ref{sec:obs}). The results are then presented in two ways: in section \\ref{sec:secular} we describe the time-averaged spectrum obtained at each epoch of observation, and then in \\ref{sec:timevar} we re-present the data as spectral time-series at 30-sec time-resolution. In section \\ref{sec:discus}, the results obtained are put into context and we review the variability uncovered in IX~Vel and V3885~Sgr and compare with the prior examples of BZ~Cam and V603~Aql. Prompted by our finding that wind activity seems to show no correlation with luminosity (thereby challenging the pure radiation-driven disk wind model), we revisit earlier-generation {\\it International Ultraviolet Explorer} (\\iue) spectroscopy of IX~Vel in order to seek further evidence of this (section \\ref{subsec:discus3}). We close with some thoughts on the next move in the quest for a fuller understanding of the driving of these accretion disk winds. ", + "conclusions": "\\label{sec:discus} In this paper we have presented high-time and wavelength resolution UV observations of the nova-like variables IX Vel and V3885 Sgr. We begin the discussion of the outcome of these observations by reviewing some of the binary parameters and other specifics of the two target systems (section~\\ref{subsec:discus1}). Next, we evaluate what has been learned from the time-series observations and compare with the already published studies of BZ~Cam and V603~Aql (section~\\ref{subsec:discus2}). We then reconsider the suitability of the line-driven disk wind model in the light of the new and, indeed, some older observations (section~\\ref{subsec:discus3}). Section 5.4 contains some comment on the way forward from here. \\subsection{IX~Vel and V3885~Sgr} \\label{subsec:discus1} First we compare and contrast IX~Vel and V3885~Sgr, as known from the literature, and as we have found them here. Sufficiently precise parallaxes exist for both binaries that we can comment on their likely relative luminosities. IX~Vel is listed in the HIPPARCOS catalogue as having a parallax of 10.38 $\\pm$ 0.98 mas, placing it at a distance of 96$^{+11}_{-8}$ pc. V3885~Sgr has a rather less certain parallax of 9.11$\\pm$1.95 mas, suggesting a distance of 110$^{+30}_{-20}$ pc. Combining these data with the impression from photometry that V3885~Sgr is typically a magnitude (i.e. close to a factor of 2) fainter at optical wavelengths, one may deduce that V3885~Sgr is anywhere between about as optically luminous as IX~Vel and half as luminous. The difference in apparent brightness is much the same at UV wavelengths, as measured in these data (cf. I1 and V3) and indeed in the older \\iue\\ and more recent \\fuse\\ data available from archive. The orbital inclination of IX~Vel would appear to be relatively settled at $i=60\\pm5\\degr$ (Beuermann \\& Thomas 1990). The situation regarding V3885~Sgr is less clear: all that can be said is that the binary is not eclipsing and that there is a difference of opinion in the literature, with Cowley et al. (1977) favouring $i\\la50\\degr$ and Haug \\& Drechsel (1995) preferring either $60\\degr$ or $70\\degr$. In the absence of better observational constraints, it would appear that to first order IX~Vel and V3885~Sgr are at similar inclinations. A problem for the analysis of our V3885~Sgr observations has been the absence of a compelling measurement of the semi-amplitude of the white-dwarf $K$-velocity, and great uncertainty in the orbital ephemeris. It is to be hoped that this can be corrected soon by means of ground-based optical spectroscopy. Against this background, it is not too surprising that the \\hst/STIS UV spectral characteristics of the two systems are not very different. In the mean of the index, the slope of the UV continuum in the two objects is about the same, fitting the power law $F_{\\lambda} \\propto \\lambda^{-2.3}$, but we do see more variation in continuum slope in IX~Vel. Tentatively we find that a redder UV continuum in IX~Vel (power law index $-2.3$) is linked to more extreme spectral evidence of mass loss -- when the wind is all but absent a bluer slope of index $-2.7$ is observed. This leads us to surmise the presence of some diffuse Balmer continuum emission due to the wind that acts to redden the spectrum slightly when the wind blows. If the same is happening in V3885~Sgr, it may be less evident because the variation in wind activity is less marked and perhaps (more speculatively) because the binary's orbit is less highly-inclined to the line of sight. Both IX Vel and V3885~Sgr show evidence of mass loss in \\lya~as well as in the often-remarked C\\IVl1549, Si\\IVl1398 and N\\Vl1240 resonant transitions. V3885~Sgr distinguishes itself from IX~Vel in also presenting with blueshifted absorption in Si\\IIIl1206 -- this is particularly evident in V3, the third epoch data wherein the wind spectral signatures were most pronounced. In general, IX Vel appears to possess the more opaque outflow. This can be seen by comparing the strength of absorption in the two datasets that display the most well-developed wind profiles: I1 and V3. In particular, the two most opaque transitions, N\\Vl1240 and C\\IVl1549 have greater equivalent widths in I1 than in V3 (see tables \\ref{tab:ixvel_vels} and \\ref{tab:v3885_vels}). Despite this evidence of a difference in typical wind optical depths, it cannot yet be determined whether this is the effect of a greater global mass loss rate or simply an orbital inclination effect. However, what can be said is that the lower wind opacity in V3885 Sgr is likely to make it the better testbed for probing links between the wind and the physics of the disk it emerges from. \\subsection{Time variability, and a comparison with the stronger UV spectroscopic variables BZ~Cam and V603~Aql} \\label{subsec:discus2} In section \\ref{sec:timevar} we have presented trailed mean-subtracted spectra of the N\\Vl1240, Si\\IVl1398 and C\\IVl1549~lines. The time resolution of these series is 30\\thinspace sec. In one sense these time series are a disappointment because the broad blueshifted absorption is revealed as being relatively steady in both binaries. There is none of the more obvious short term variability described in the UV spectra of V603~Aql (Prinja et al. 2000a) and BZ Cam (Prinja et al. 2000b). A re-examination of IX Vel difference spectra at 10-sec time-resolution has revealed no variability unseen in the 30-sec time-binning. It would appear then, on the limited basis of a sample of 4 objects, that (i) systems selected for examination at higher time resolution, because they are already known to be significantly variable on the longer timescale of tens of minutes, continue to reward with variability on the tens of seconds timescale (BZ Cam, V603~Aql), (ii) systems not known to be markedly variable on the longer timescale remain as such on the shorter (IX Vel, V3885~Sgr). A possible point of distinction between these two pairs of objects is that the variable objects are also mooted as very low orbital inclination binaries ($i < 40^{\\circ}$). This difference permits the hypothesis that the variability is due to the line-of-sight sampling some fluctuation in the angling of the innermost disk-wind streamlines. This assumes that there is something like a central conical cavity in the outflow as sketched by Shlosman \\& Vitello (1993) and by Knigge, Woods \\& Drew (1995) and substantiated by the numerical models published by e.g. PSD. By contrast, at higher inclination the line-of-sight will be well and truly embedded in the body of the approaching outflow at all times. Clearly it would be of benefit to be more sure of these binaries' orbital inclinations. There is another very striking point of contrast to note between these two pairs of objects. In the UV spectra of IX~Vel and V3885~Sgr, there is very little redshifted line emission partnering the broad blueshifted absorption -- only C\\IVl1549 typically shows any at a significant level. This is quite common among high state non-magnetic CVs. Yet in BZ Cam and V603~Aql there is very strong redshifted line emission. This difference could be said to run right against expectation if BZ~Cam and V603~Aql really are much more nearly face-on binaries, given that strong line emission is generally associated with eclipsing systems ($i\\ga70\\degr$). In this context it is helpful to remember that exercises in the simulation of wind-formed UV line profiles (e.g. Drew (1987), Shlosman \\& Vitello (1993), Knigge et al. (1995)) together with the discovery of orbital-phase linked line profile variations have led to the recognition that a component of the detected UV line emission need not have a wind origin (see discussion of this in Drew, 1997). Given this, the implication of the stronger redshifted line emission in BZ~Cam and V603~Aql is either that these binaries are `the same' as IX~Vel and V3885~Sgr and somehow the line emission is occulted at higher orbital inclinations or there is an intrinsic physical difference between these two pairs of objects that expresses itself via the observed strength of the UV line emission. In both IX~Vel and V3885~Sgr there is evidence that narrow absorption features, sometimes seen superposed on broad blueshifted absorption, exhibit an organised motion that may prove to follow that of the accreting WD. At a mean blueshift of $\\sim-900$ \\kms, these features are present only in the data showing the most pronounced P~Cygni behaviour. In the I1 and V3 timeseries of the mean-subtracted line profiles, the red absorption edge close to line centre approximately follows the motion of the narrow features -- as makes sense if both are due to the wind and hence follow the binary motion of the white dwarf. If this interpretation is correct, and there is no azimuthal dependence of the line-of-sight velocity of these features, the detected velocity changes provide a constraint on the orbital phase observed and the white-dwarf $K$-velocity. It is likely that the manner in which these narrow absorption features vary gives a clue as to their origin. On noticing them in \\iue\\ high-resolution data on IX Vel, Mauche (1991) reported three categories of explanation for similar features seen in O-star spectra as summarised by Henrichs (1988). The first concept is that an absorption dip may be caused by a plateau in the outflow velocity law which allows a larger absorbing column to build up over a small change in line-of-sight velocity. The second is that an absorber density enhancement due to e.g. an ionization effect may have the same consequence. Lastly, one may appeal to non-monotonicity due to e.g. shocks sweeping through the flow as the cause of narrow dips. Only the velocity law plateau concept warrants further examination in the context of these disk winds. This is because there are obvious problems with the other two. First, the density enhancement idea sits uneasily with the observed presence of these dips in transitions sampling a wide range of ionization states together with their absence from a subordinate line (C\\IIIl1175) that should be favoured at an enhanced density. Secondly, the difficulty with the shock-swept, or otherwise unstable, disk wind class of explanation is the persistence and near stationarity of the dips at a velocity well below the terminal value. Table \\ref{tab:iue_vels} list those \\iue\\ datasets in which there are clearly narrow features superposed on the C\\IVl1549 line, along with the locations of these narrow features (see also Mauche 1991, Prinja \\& Rosen 1995). The similarity between the dip blueshifts in IX~Vel and V3885~Sgr, together with their earlier appearance in \\iue\\ observations, offers the prospect of an origin rooted in some basic property common to both binaries. \\begin{table} \\caption{Location of narrow features found superimposed on the C\\IVl1549 doublet in high resolution \\iue\\ data of IX Vel. Velocities are given in \\kms.} \\label{tab:iue_vels} \\begin{tabular}{@{}lll} \\hline \\iue\\ dataset\t&C\\IVl1548.2\t&C\\IVl1550.7\t\\\\ \\hline SWP22353\t&$-1730\\pm70$,$-850\\pm70$&$-960\\pm60$\t\t\\\\ SWP25449\t&$-1000\\pm90$\t\t&$-980\\pm80$\t\t\\\\ SWP29618\t&$-1040\\pm70$\t\t&$-1010\\pm70$\t\t\\\\ \\end{tabular} \\end{table} It should be mentioned that similar narrow absorption dips are seen in association with O\\VIl1035 blueshifted absorption in outburst spectra of U~Gem \\cite{2001froning}. These data are complementary to those presented here in that the orbital phase coverage is close to complete, while the spectral resolution and S/N ratio is reduced. The greater duration of the FUSE observations reveals that the O\\VI\\ dips vary about their mean blueshift of $\\sim 500$~km s$^{-1}$ by no more than a few tens of kilometres per second. The more marked variation is in equivalent width -- the dips are seen to fade and reappear in a somewhat erratic fashion. In one epoch of observation of V3885~Sgr, V2, the mean-subtracted timeseries reveals a quite subtle and gradual change of a different character: the seeming weakening of the wind signature in e.g. the C\\IV~line profile during the second 1000\\thinspace sec of the observation. This might be interpreted as a temporary calming of the outflow. The change, at the same time, in the appearance of the 1300\\thinspace\\AA\\ absorption feature (figure~\\ref{fig:v3885_lines}) points to a different explanation. This typically broad shallow depression is thought of as originating in the disk atmosphere. The fact that it develops superposed narrow absorption dips at close to the rest wavelengths of the Si\\III~multiplet, the dominant contributor, strongly hints at a change in the visibility of the inner disk -- and along with it, the disk wind. This is not exceptional behaviour, in that similarly-narrow unshifted absorption dips are the norm for U~Gem in outburst \\cite{2001froning}, and have been observed at a range of orbital phases in the UV spectra of UX~UMa and RW~Tri (Mason, Drew \\& Knigge 1997). That the narrow absorption is fleeting in V3885~Sgr may be a hint of a somewhat lower, but not too low, orbital inclination. Finally, we do not find frequent signatures of wind inhomogeneity on or near the dynamical timescale in our data for either IX~Vel or V3885~Sgr - just one spell of variability, the $\\sim350$\\thinspace sec `wobble' seen in the V3 dataset is all there is to report. This indicates that structure extending over a significant fraction of the projected area of the UV continuum source is rare. If the inhomogeneity described by PSD exists, it must be finer in scale. One way of achieving this in a model may be to allow the flow to be fully 3-dimensional. Specifically, the derived geometry and dynamics of a 3-D wind is likely to differ from those of the axisymmetric 2-D wind modelled by PSD. For example, the knots discovered by PSD correspond to rings in 3-D with axisymmetry imposed, but in a full 3-D simulation they may be replaced by spirals or even a foam of very small blobs. If there really is no significant inhomogeneity, then the radiation-driven wind models of PSD would demand that around a half or more of the luminosity is radiated quasi-spherically, or that flow co-rotation is enforced by e.g. a strong, large-scale, ordered magnetic field threading the disk. However, before such adjustments are considered, there is a more basic aspect of the radiation-driven wind model to review. \\subsection{Implications for the mass loss mechanism} \\label{subsec:discus3} If radiation pressure powers the mass loss from IX~Vel and from V3885~Sgr, there are good reasons to expect a strong scaling of the rate of mass loss to bolometric luminosity. Indeed, using the plot of effective Eddington number ($\\Gamma$, the ratio of bolometric luminosity to the Eddington limit) against mass loss rate that was given in Drew \\& Proga (2000), it is possible to estimate how strong this dependence should be. On adopting the mass accretion rate and WD parameters derived for IX~Vel by Beuermann \\& Thomas (1990), we would expect $\\Gamma$ to be about $4\\times10^{-4}$. For V3885~Sgr, this quantity is unlikely to be greater and is probably smaller. This places both systems in the regime where an exponential dependence is expected. Even if it were the case that $\\Gamma$ were high enough to yield a gentler dependence, this would still be at least as steep as $\\mdot_{\\rm{w}} \\propto \\lbol^2$. So, unless there is a significant countermanding effect due to shifts in wind ionization, we should expect to see a strong positive correlation between the strength of wind features in the ultraviolet and the ultraviolet continuum level -- given that the latter is expected to scale very nearly linearly with mass accretion rate for fixed WD parameters. The time-averaged spectra presented in section 3 do not sit well with this expectation. In V3885~Sgr the weakest wind signatures are seen when the ultraviolet continuum is brightest (V2), and strongest when the UV continuum is 20 to 30 per cent fainter (V3). In our IX Vel data (I3), the wind all but disappears when the UV continuum is at a level intermediate between the levels observed at the other two epochs. This absence of trend is seen in a setting where the evidence from the UV spectral lines is that ionization shifts are not significant and abundant ion stages are represented (see section 3.1). This gives us confidence that there is a direct (if uncalibrated) mapping from the mean strength of blueshifted absorption onto mass loss rate, from the lines that we observe with \\hst. Caution must be exercised regarding the significance of IX Vel's apparent UV flux changes because these data were obtained through an aperture small enough that the STIS flux throughput could have been different at each epoch. However, it is very unlikely that we are misled in the case of V3885~Sgr since the data were obtained through the larger 0.2 arcsec$^2$ aperture that has been characterised as yielding UV fluxes to with 10 per cent of their photometric values. \\begin{figure*} \\vspace*{12cm} \\special{psfile=fig13.eps hscale=82 vscale=62 voffset=-100 hoffset=10} \\caption{I1 and I3 rebinned to $\\Delta\\lambda=1.8$\\thinspace\\AA\\ and overplotted on an \\iue\\ low-dispersion spectrum. Top frame - I1 plotted over data set SWP19765 . Bottom frame - I3 plotted over data set SWP42994. The \\hst\\ I1 and I3 spectra have both been multiplied by constants in order to rescale their flux levels to those of the \\iue\\ spectra with which they are compared.} \\label{fig:ix_iue} \\end{figure*} Since the failure to find a correlation between wind features and continuum level is inconsistent with our expectation, it is important to see whether independent evidece of the effect can be found in the \\iue\\ database. We have done this for the case of IX Vel as there is a particularly extensive set of \\iue\\ spectra to draw upon (12 at high resolution and 28 at low resolution). First of all, we note that the \\iue\\ recorded flux level in the 1260--1270\\thinspace\\AA\\ range varies from $9.1\\times 10^{-12}$ to $13.6 \\times 10^{-12}$ erg s$^{-1}$ cm$^{-2}$\\thinspace\\AA$^{-1}$. The \\hst\\ data at epochs I1 and I3 (table 3) are consistent with this, while I2 was a time of seemingly unusually low continuum flux. Given that the $V$ band flux is known to vary secularly by up to $\\sim$0.9 mags, the factor of 2 drop between I1 and I2 is not so large it must be instrumental in origin. Secondly, we have rebinned the weak-wind I3 spectrum to match the low resolution mode of \\iue\\ in order to determine if such a state was also recorded by \\iue. We find that three of the 28 `LORES' spectra (SWPs 19765, 19766 and 26081) resemble the rebinned I3 data: to illustrate this, in figure \\ref{fig:ix_iue} an overplot of the rebinning on SWP 19765 is presented. If the I1 data are similarly degraded in spectral resolution, these too can be matched to examples of \\iue\\ spectra (see also figure \\ref{fig:ix_iue}). These simple comparisons allow the conclusion that the \\hst\\ observations describe the same basic phenomenology as \\iue\\ (but more fully, of course). \\begin{figure} \\vspace{16cm} \\special{psfile=fig14a.eps hscale=40 vscale=40 voffset=170 hoffset=0} \\special{psfile=fig14b.eps hscale=40 vscale=40 voffset=-60 hoffset=0} \\caption{Plot of C\\IV$\\lambda1549$ absorption equivalent width against, a) continuum flux in the 1260--1270\\thinspace\\AA\\ range (above) and b) spectral index (below) for IX Vel \\iue\\ archive data. Solid triangles represent LOWRES data and open triangles represent HIRES data. An example of the estimated typical error bar is given in the top left-hand corner of each plot. Pearson correlation coefficients for a linear fit to these data are given in table \\ref{tab:pearson}.} \\label{fig:ewflux} \\end{figure} One of the admirable features of the final archive of \\iue\\ data is its photometric consistency. This allows us to carry out the experiment on IX~Vel archive data of evaluating the strength of correlation between wind activity, as signalled by the equivalent width of C\\IVl1549 blueshifted absorption, with (i) UV flux level, again measured over the narrow band, 1260--1270\\thinspace\\AA, (ii) UV continuum spectral index (fitting the continuum to a power law using much the same method as adopted for the \\hst\\ summed spectra). These pairs of quantities as measured from \\iue\\ archive data are plotted in figure \\ref{fig:ewflux}. In addition we have formally evaluated the Pearson correlation coefficients between these variable pairings, and also between the UV flux level and spectral index The results obtained are shown in table \\ref{tab:pearson} . In support of what we have seen in the \\hst\\ data, there seems to be at best scatter, or perhaps even a weak anti-correlation, between the C{\\sc iv} absorption equivalent width and UV brightness (figure \\ref{fig:ewflux}a). By contrast, there does seem to be a correlation between the C\\IV\\ line and UV spectral index in the same sense as picked out earlier -- i.e. the slope of the spectrum reddens as the wind activity increases. \\begin{table} \\caption{Pearson's correlation coefficient, r, for the relationship between continuum spectral index, flux and C\\IV\\ blueshifted absorption equivalent width in the \\iue\\ archive IX Vel data. $P(r)$ is the probability of a random distribution having a Pearson coefficient $\\ge|r|$} \\label{tab:pearson} \\begin{tabular}{@{}lcc} \\hline Variables\t\t&$R$\t&$P(r)$\t\\\\ \\hline Spectral Index vs. Flux\t&$-0.249$\t&$0.127$\\\\ C\\IV\\ EW vs. Flux\t&$-0.270$\t&$0.096$\\\\ C\\IV\\ EW vs. Spectral Index\t&$0.622$\t&$2.36\\times10^{-5}$\\\\ \\end{tabular} \\end{table} Circumstantially, if not absolutely conclusively, it would appear that there is not the expected strong positive correlation between apparent mass loss and luminosity. A key prediction of the radiatively-driven disk wind model is flouted. The flow timescale in cataclysmic binaries like IX~Vel and V3885~Sgr is likely to be a few tens of seconds (assuming a Castor \\& Lamers type velocity law from 40 to 5000\\kms integrated over a distance $10\\,R_{\\rm{WD}}$). This is sufficiently short that the wind dynamical configuration is able to adjust continuously to the changing radiant flux as long as this does not vary markedly on an even shorter timescale. In neither object is there evidence of this either in our data or, historically, from observations of flickering (e.g. Williams \\& Hiltner 1984, Cowley et al 1977). Hence, it is not an option to invoke a concept of e.g. hysteresis to explain the absence of the expected steep relation between luminosity and mass loss rate. We conclude that neither IX~Vel nor V3885~Sgr knows they have to obey such a relation! \\subsection{Final Remarks} \\label{subsec:discus4} Now that quantitative models of radiation-driven disk winds have appeared, and there are high quality data to match them. it can be seen more clearly that a factor other than line-driving is much more likely to be decisive in powering these outflows. Our findings add to the doubts over the years that line-driving cannot sustain the mass loss rates observation seems to imply (see Mauche \\& Raymond 2000, Drew \\& Proga 2000). The de facto exclusion of a class of models is progress. It is now time to examine the alternatives with renewed intent. MHD driving has long been viewed as a possible contributor or sole cause of mass loss (Cannizzo \\& Pudritz 1988, see also comments in Drew 1991). This option has become more attractive again in recent years as the magneto-rotational instability has emerged as a favoured angular momentum transport mechanism for CV disks (Balbus \\& Hawley 1998, Tout 2000). Perhaps the most promising alternative to a pure line-driven disk wind model is a hybrid model where a wind is driven from the disk by some combination of MHD and line forces. Calculations combining the two driving mechanisms rigorously do not exist yet. However we can refer to some schematic predictions of radiation-driven disk wind models where the main role of an ordered magnetic field is to impose angular velocity, rather than angular momentum, conservation (Proga 2000). For $L_{\\rm{D}} \\ga 0.001 L_{\\rm{E}}$, imposing corotation on a line-driven disk wind model increases the wind mass loss rate: the smaller the angle between the poloidal streamlines and the disk midplane the larger mass loss rate (cf. Blandford \\& Payne 1982). For $L_{\\rm{D}}>$ a few times $0.001L_{\\rm{E}}$, $\\mdot_{\\rm{w}}$ is as in a pure line-driven disk wind model. In terms of the $\\mdot_{\\rm{w}}$ vs $L_{\\rm{D}}$ relationship, it means that for a relatively low disk luminosity (i.e. $L_{\\rm{D}} \\ga 0.001 L_{\\rm{E}}$), the mass loss rate can increase or decrease depending on the streamline geometry required by prevailing magnetic fields. This control of $\\mdot_{\\rm{w}}$, independent of the radiant luminosity, can potentially explain our observations because these sample the luminosity range where this additional dependence is expected. As noted already, enforcement of flow streamlines by a large scale magnetic field would also be consistent with the near absence of detected wind inhomogeneity. On the other hand, the fact that a radiation-driven wind can be produced only when $L_{\\rm{D}} \\ga 0.001 L_{\\rm{E}}$, has a basic appeal in explaining the behaviour of other CVs such as dwarf novae (DN). The UV observations of DN through outburst obtained by \\iue\\ have provided evidence of mass loss signatures tracking luminosity. For example, it has been noted during declines of DN from maximum light that a decay of the UV continuum by a factor of $\\ga 2$ can be accompanied by a near disappearance of the blueshifted absorption so prominent at maximum (Woods \\& Drew 1990; Woods et al. 1992). If this correlation proves to be indirect rather than causal, we have work to do to uncover the controlling physics behind it. The way forward is to identify for further development one of the many variants of MHD disk wind model already explored in the literature that can provide a convincing explanation, either with or without the assistance of line-driving, of the continuing puzzle of these disk winds. Once again the accretion disk winds in CVs have emerged as particularly open and rewarding laboratories for studying disk mass loss. An intriguing issue for the future is the existence or not of flow collimation of CV disk winds if even CV disk winds are MHD powered (see Knigge \\& Livio 1998). As these problems are tackled we can also hope to obtain greater insight into the nature of disk winds in other astrophysical settings." + }, + "0112/astro-ph0112338_arXiv.txt": { + "abstract": "An analytical expression is presented that allows gas-to-dust elemental depletions to be estimated in interstellar environments of different types, including Damped Ly\\,$\\alpha$ systems, by scaling an arbitrary depletion pattern chosen as a reference. As an improvement on previous work, the scaling relation allows the dust chemical composition to vary and includes a set of parameters which describe how sensitive the dust composition is to changes in both the dust-to-metals ratio and the composition of the medium. These parameters can be estimated empirically from studies of Galactic and extragalactic depletion patterns. The scaling law is able to fit all the typical depletion patterns of the Milky Way ISM (\\emph{cold disk}, \\emph{warm disk}, and \\emph{warm halo}) with a single set of parameters, by only varying the dust-to-metals ratio. The dependence of the scaling law on the abundances of the medium has been tested using interstellar observations of the Small Magellanic Cloud (SMC), for which peculiar depletion patterns have been reported in literature. The scaling law is able to fit these depletion patterns assuming that the SMC relative abundances are slightly non solar. ", + "introduction": "It is well known that elemental abundances measured in the interstellar gas of the solar vicinity are generally depleted with respect to the solar values (Morton 1974; Jenkins, Savage \\& Spitzer 1986). The commonly accepted interpretation of this depletion effect is that a fraction of these elements is not detected in the gas because it is locked into dust grains. Elemental depletions are estimated by comparing the abundances measured in the gas with the abundances of the medium (gas plus dust). Intrinsic solar abundances are usually adopted in local interstellar studies, even though the interstellar abundance standard is not yet firmly established (Savage \\& Sembach 1996, hereafter SS96; Sofia \\& Meyer 2001). Different elements may have very different values of depletion, the so called \"refractory\" elements being almost completely in dust form. In addition, elemental depletions can vary significantly among different lines of sight. In spite of this complex behavior, interstellar regions with similar physical conditions are characterized by similar depletions (Spitzer 1985; Jenkins et al. 1986). In the Galactic disk the cold gas has higher depletions than the warm gas; in the halo the warm gas shows even lower depletions (SS96). Elemental depletions constitute an important piece of information for solving the complex puzzle of the origin and evolution of interstellar dust (Tielens \\& Allamandola 1987 and refs. therein). Studies of depletions in the Galaxy and in the Magellanic Clouds (MCs) have received new impulse (Roth \\& Blades 1995; Welty et al. 1997, 2001) owing to their importance in the interpretation of abundance patterns measured in the QSO absorbers of highest \\ion{H}{1} column density, namely the Damped Ly $\\alpha$ systems (DLAs). These systems originate in intervening galaxies or protogalaxies and their abundances yield unique information on galactic chemical evolution at different redshifts (Lu et al. 1996; Prochaska \\& Wolfe 1999; Molaro et al. 2000). The evidence that dust depletion may significantly affect the abundances of DLAs is quite compelling (Pettini et al. 1994; Lauroesch et al. 1996; Hou et al. 2001; Prochaska \\& Wolfe 2002). In order to cope with the problem of depletion in DLAs some authors have concentrated on studies of non-refractory elements (Pettini et al. 1997; Centuri\\`on et al. 2000; Vladilo et al. 2000) and of systems with small dust content (Pettini et al. 2000, Molaro et al. 2000). At the same time, formalisms have been developed to quantify the effect of dust depletions on DLA abundances (Kulkarni et al. 1997; Vladilo 1998; Savaglio, Panagia \\& Stiavelli 2000). These formalisms assume that the dust in DLAs is similar to Galactic interstellar dust. Depletions in DLAs have been modeled scaling Galactic depletions by allowing variations of the dust-to-metals ratios, but not of the dust chemical composition. In addition, the fact that the dust composition may be affected by variations of the intrinsic abundances of the intervening galaxies has not been considered. In this work a scaling law of interstellar depletions is derived that allows the dust chemical composition to vary according to changes in the dust-to-metals ratio and/or to changes in the overall abundances of the medium (Section 2). Rather than assuming {\\it ad hoc} types of dust, a set of parameters is introduced to describe how sensitive the dust composition is to changes in the physical and chemical properties of the medium. In Sections 3 and 4 it is shown that such parameters can be estimated from Galactic and extragalactic interstellar data. The conclusions are summarized in Section 5. ", + "conclusions": "An analytical expression has been derived that allows dust depletions to be estimated in interstellar environments characterized by a wide range of chemical and physical properties. In practice, the depletions are estimated as a function of the dust-to-metals ratio, $r$, and of the relative abundance of the element X in the medium, $a_\\mathrm{x}$, by scaling a depletion pattern chosen as a reference. The scaling law, given in Eq. (\\ref{ScalingLawGeneral}), has been derived making no assumption on the mechanisms of dust formation, accretion or destruction. No hypothesis has been made on the extinction properties of the dust. The functional form (\\ref{ScalingLawGeneral}) is always valid for infinitesimal changes $\\delta r$ and $\\delta a_\\mathrm{x}$. For finite changes $\\Delta r$ and $\\Delta a_\\mathrm{x}$ the expression is valid if the parameters $\\eta_\\mathrm{x}$ and $\\varepsilon_\\mathrm{x}$ are constant. These parameters are essentially derivatives of the relative abundance of the element X in the dust, $p_\\mathrm{x}$. In practice, the parameters $\\eta_\\mathrm{x}$ indicate how the chemical composition of the dust is affected by changes of the dust-to-metals ratio $r$. The parameters $\\varepsilon_\\mathrm{x}$ indicate how the composition of the dust is affected by changes in the composition of the medium. The scaling law can be applied to interstellar clouds of our Galaxy and of external galaxies with solar or non-solar abundances --- including Damped Ly $\\alpha$ systems --- once the sets of parameters $\\eta_\\mathrm{x}$ and $\\varepsilon_\\mathrm{x}$ are determined. The assumption required to apply the scaling law to external galaxies is that the dust chemical composition is a function of $r$ and $a_\\mathrm{x}$ valid in any type of interstellar medium. This assumption is much more realistic than assumptions previously adopted in studies of DLA systems, in which depletion patterns were estimated using specific types of Galactic dust. The parameters $\\eta_\\mathrm{x}$ and $\\varepsilon_\\mathrm{x}$ can be determined empirically by comparing observed depletion patterns in interstellar media with different physical and chemical properties. We have used the typical Milky Way depletion patterns to estimate the parameters $\\eta_\\mathrm{x}$ in environments with variable $r$ but constant (solar) abundances. The resulting values of $\\eta_\\mathrm{x}$ are approximately constant for most elements in a wide range of physical conditions, supporting the validity of the scaling law (\\ref{ScalingLawGeneral}), at least as far as the dependence on $r$ is concerned. With the derived set of $\\eta_\\mathrm{x}$ parameters, the scaling law is able to simultaneously fit all the typical Milky Way depletion patterns by only varying $r$. The parameters $\\varepsilon_\\mathrm{x}$ can be determined empirically by studying interstellar depletions of galaxies with known abundances different from the solar ones. We have applied this technique to the SMC, for which the intrinsic abundances are constrained from stellar studies and depletion data are available for three lines of sight. The observed SMC depletions can be modeled with the scaling law if the SMC abundance ratios are slightly non solar. The required departures from solar ratios are consistent with the results of SMC stellar studies and with expectations based on Galactic stellar abundances of stars having the same metallicity of the SMC. The agreement between observed and predicted depletions is obtained at values $\\varepsilon_\\mathrm{x} \\simeq 1$ or somewhat below unity. With this choice of parameters the scaling law is able to simultaneously fit the three available SMC depletion patterns, including the anomalous pattern observed toward Sk 155. Values $\\varepsilon_\\mathrm{x} > 1$ are excluded by the present analysis. The uncertainty of the SMC stellar abundances prevent firmer conclusions or more constrained $\\varepsilon_\\mathrm{x}$ values. The capability of the scaling relation presented here to match Galactic and SMC depletion patterns is rather encouraging. To probe the general validity of the scaling law it is desirable to measure depletions in a larger number of extragalactic lines of sight. However, these studies will yield stringent constraints only if the intrinsic abundances of the external galaxies, mostly based on stellar observations, will be determined with accuracy. A major step in this direction can be made in the next few years by observing a relatively large number of stars of the Magellanic Clouds at high spectral resolution. These time-consuming efforts will be rewarded by the possibility of modeling interstellar depletions and deriving dust-corrected abundances in high redshift galaxies observed as QSO absorption systems. An implementation of the scaling relation for deriving dust-corrected abundances in Damped Ly\\,$\\alpha$ systems will be presented in a separate paper." + }, + "0112/astro-ph0112174_arXiv.txt": { + "abstract": "We have imaged the emission from the near-infrared v=1--0 S(1), 1--0 S(7), 2--1 S(1) and 6--4 O(3) lines of molecular hydrogen in the N-- and SW--Bars of M\\,17, together with the hydrogen Br$\\gamma$ and Br10 lines. This includes the first emission line image ever to be obtained of a line from the highly excited v=6 level of H$_2$. In both Bars, the H$_2$ emission is generally distributed in clumps along filamentary features. The 1--0 S(1) and 2--1 S(1) images have similar morphologies. Together with their relative line ratios, this supports a fluorescent origin for their emission, within a photodissociation region. The SW--Bar contains a clumpy medium, but in the N--Bar the density is roughly constant. The 1--0 S(7) line image is also similar to the 1--0 S(1) image, but the 6--4 O(3) image is significantly different to it. Since the emission wavelengths of these two lines are similar (1.748 to 1.733\\mic), this cannot be due to differential extinction between the v=6 and the v=1 lines. We attribute the difference to the pumping of newly formed H$_2$ into the v=6, or to a nearby, level. However, this also requires either a time-dependent photodissociation region (where molecule formation does not balance dissociation), rather than it to be in steady-state, and/or for the formation spectrum to vary with position in the source. If this interpretation of formation pumping of molecular hydrogen is correct, it is the first clear signature from this process to be seen. ", + "introduction": "\\subsection{Messier 17} Messier 17 is one of the nearest regions of massive star formation to us, situated $\\sim 1.3$\\,kpc away (Hanson, Howarth \\& Conti 1997). It is particularly notable for the bright \\hii\\ region which attracted Messier's attention (and also known as the Omega Nebula). Within it lies an obscured cluster of young stars. Thirteen OB stars have been spectroscopically identified (Hanson et al.\\ 1997), all with between 4 and 15 magnitudes of extinction at optical wavelengths and estimated to be $\\sim 10^6$\\,years old. The \\hii\\ region is adjacent to a molecular cloud complex, in particular the core M\\,17 SW, which contains more than $\\rm 10^4 M_{\\sun}$ of material (Lada 1976, Thronson \\& Lada 1983). Intense far--IR emission arises from the core, with a luminosity of $\\rm \\sim 6 \\times 10^6 L_{\\sun}$ (Harper et al.\\ 1976, Wilson et al.\\ 1979). While this originates from dust exposed to far--UV radiation, the core appears to be externally heated, by the stars that excite the \\hii\\ region (Gatley et al.\\ 1979), rather than internally. It forms a photodissociation region (PDR, Tielens \\& Hollenbach 1985). The method of excitation of the hydrogen molecules in it are the subject of this paper. \\subsection{Molecular Hydrogen Observations in M\\,17} M\\,17 was first mapped in H$_2$ through the 2.12\\mic\\ 1--0 S(1) line by Meadows (1986) and Gatley \\& Kaifu (1987). Their maps reveal two obvious regions of emission, a bar running across the north, and another running south to west. They are known as the `Northern Bar' and the `South Western Bar', respectively. They closely follow two emission ridges seen in radio continuum from 1.3 to 21\\,cm (Felli et al.\\ 1984). These bars are PDRs, interface regions between the ionized and the molecular gas. The Northern Bar also follows an optically visible bar of the \\hii\\ region, suggesting that the extinction to it is relatively small. The South Western Bar, however, is optically obscured. PDR modelling of far--IR and sub--mm line emission features from the SW--Bar (Burton, Hollenbach \\& Tielens 1990, Meixner et al.\\ 1992) indicates that it is a clumpy region, containing gas with a mix of densities, ranging from $10^3$ to $10^7$\\,\\cmt, exposed to a far--UV radiation field of $\\sim 10^4$ times the average interstellar radiation field. Low spatial resolution (20 arcsec beam) near--IR spectroscopy of the 1--0 and 2--1 S(1) H$_2$ lines (Tanaka et al.\\ 1989) in the N--Bar found a line ratio of $\\sim 2$. This is the ratio expected if the excitation is via a fluorescent cascade through the vibrational-rotational levels of the ground electronic state (Black \\& Dalgarno 1976). The molecule is first excited by far--UV photons to an excited electronic level, from which it decays. Chrysostomou et al.\\ (1992, 1993) imaged the N--Bar with arcsecond spatial resolution in the H$_2$ 1--0 S(1) line, as well as in the 3.29\\mic\\ PAH emission feature and the hydrogen Br$\\gamma$ line at 2.17\\mic. They also obtained long slit spectroscopy across the N--Bar from 2.0--2.5\\mic, measuring 16 lines of H$_2$, as well as several \\hii\\ region emission lines. The spatial coincidence found for the PAH and H$_2$ emission features confirmed the fluorescent origin, and also suggested that most of the gas had density $> 10^5$\\cmt. They found the 1--0/2--1 S(1) ratio to be $\\sim 3$, and to be constant along the 60\\,arcsec long slit. This not only suggests that the H$_2$ is slightly thermalised (i.e.\\ the density is about the critical density, $\\sim 10^5$\\cmt), but that the physical conditions along the N--Bar are constant. \\subsection{Molecular Hydrogen Excitation in PDRs} In dense PDRs, with a strong far--UV radiation field, self-shielding can bring the H/H$_2$ dissociation front to optical depths $A_v < 1$\\,mag from the ionized surface of the cloud. Here the temperature can reach 1,000--2,000\\,K, allowing the v=1 level of H$_2$ to be thermalised. H$_2$ line ratios, for instance the 1--0/2--1 S(1) ratio, then increase from the pure fluorescent value of 2, towards the thermal value of 10 observed in shocks and in gas at $\\sim 2000$\\,K (see Sternberg \\& Dalgarno 1989, Burton, Hollenbach \\& Tielens 1990). The actual line ratio depends on the density of the gas. If this is not constant, then the ratio will likely vary with position in a source. However calculating its value depends on an accurate knowledge of the collisional excitation rate coefficients for H$_2$, which are poorly known. Hence the quantitative application of ``collisional fluorescence'' models to the interpretation of H$_2$ lines ratios in PDRs requires some care. A component of the fluorescent emission from H$_2$ must arise due to the formation of the hydrogen molecules themselves. For photo-excitation of H$_2$ molecules by far--UV photons, approximately 15\\% lead to dissociation. However the hydrogen molecules can be reformed, through the recombination of two hydrogen atoms on the surfaces of dust grains (Hollenbach \\& Salpeter 1971). In the steady state, the rate of formation of H$_2$ molecules balances its rate of photodissociation. The 4.5\\,eV bond energy of H$_2$ is released on formation. Some of this is used in escaping from the surface of the dust grain, some goes into translational kinetic energy in the molecule, and the rest goes into internal energy. The new molecules can then emerge in a excited vibrational-rotational state, from which radiative decay will occur. Thus a formation component to the H$_2$ spectrum is expected. However, the signature of this spectrum is unclear. As Black \\& Dalgarno (1976) first realised, when calculating the infrared spectrum of fluorescent H$_2$, the contribution from molecule formation may be important. In their original model they assumed equipartition of the binding energy released on formation. They divided this equally between the internal energy of the molecule, its translational energy on escape from the grain surface, and the internal energy imparted to the grain lattice. They further assumed that the 1.5\\,eV provided as internal energy to the molecule was spread with a Boltzmann distribution through the vibrational-rotational levels. However, other assumptions are possible. For instance, Hunter \\& Watson (1978) argue that H$_2$ molecules are released in rotationally hot, vibrationally cold states (i.e.\\ high--J $\\geq 7$, low--v). A model by Duley \\& Williams (1986), on the other hand, although agreeing that the molecules should form hot, argued that they would appear in just the opposite combination of states (i.e.\\ low--J, high--v, with v=6 likely). Le Bourlot et al.\\ (1995) have investigated the infrared spectrum produced for pure formation pumping of H$_2$, under a variety of formation models. For instance, as might be expected, the intensities of v=6 lines are found to be considerably greater under the Duley \\& Williams (1986) model than in other models. However, since this model does not include the fluorescent cascade component to the emission, which will generally dominate the intensity of the H$_2$ lines, it is hard to use it to undertake a quantitative comparison with data. \\subsection{Previous Reports of the Detection of Formation Pumping for H$_2$} There have been few reports made of the signature for H$_2$ formation being observed in a spectrum. Wagenblast (1992) considers UV absorption spectra of three nearby diffuse clouds, along lines of sight to background stars. He finds that the populations in the excited, pure-rotational levels, v=0, J=5, 6 \\& 7, along the sight lines, cannot be produced by UV or thermal excitation. Assuming that molecule formation takes place in two adjacent rotational levels within a particular vibrational state (one for ortho-H$_2$, or odd--J, and the other for para-H$_2$, or even--J), he calculates the possible pairings that could account for the observed ratios of the three lines. He found that H$_2$ would need to be formed in a rotationally hot state ($\\rm J \\geq 7$), but with a range of vibrational states possible, up to v=11. Federman et al.\\ (1995) explore this further with additional absorption measurements from the v=3 level in the source $\\zeta$\\,Oph, though do not find any fit to the data particularly satisfactory. Mouri \\& Taniguchi (1995) consider the 1.5--2.5\\mic\\ spectrum of the starburst galaxy NGC 6240, where a number of H$_2$ lines are evident. Based on the relative strengths of the 1--0 S(7) and S(9) lines, compared to the 1--0 S(1) line, they argue that formation pumping, via associative detachment of H and H$^-$ to form H$_2$, provides an important contribution to the line intensity. However the data on which this is based have low spectral resolution, and suffers from blending. Moreover, the data are fit with a model which contains several components to the 1--0 S(1) line intensity; formation (10\\%), fluorescence (20\\%) and thermal excitation (70\\%). In addition, the v=6--4 Q(1) line at 1.64\\mic, observed by Elston \\& Maloney (1990) in this source, and arising from the same level as the 6--4 O(3) line which we report on in this paper, is weak compared to the 1--0 S(1) line, less than 10\\% its intensity. It has low signal to noise in the spectrum. Hence, claims for a formation signature in the data need to be regarded with caution. Measurements of over 30 high-excitation H$_2$ lines in the reflection nebula NGC~2023 by Burton et al.\\ (1992) lead these authors to speculate on whether there was a formation component to the emission from several lines from the v=4 level. This was based on an excess in the level column density distribution for these lines compared to expectations for pure fluorescent emission (and see also McCartney et al.\\ 1999 for an extension to v=6 in this source, for which the excess may still be apparent). However, the signal to noise is not sufficient to be certain that the excess is real, and moreover, might be mistaken for ortho-to-para ratio variations between lines in different vibrational levels. This paper presents images of emission lines from the v=1, 2 and 6 levels in M\\,17, as well as in the hydrogen Br $\\gamma$ line. This includes the first map ever obtained of an H$_2$ emission line from the v=6 level. While the bulk of the H$_2$ emission is clearly fluorescent in origin, we argue, based on the different morphology for the v=6--4 O(3) line to the lower excitation lines, that formation pumping of H$_2$ provides a significant component to its flux. ", + "conclusions": "\\subsection{PDR Emission and Gas Density} As has been discussed in several previous papers, the molecular hydrogen line emission in M\\,17 is dominated by UV fluorescence (e.g.\\ Tanaka et al.\\ 1989, Chrysostomou et al.\\ 1992, Chrysostomou et al.\\ 1993). This is evident through the morphology, its proximity to the ionization front of an \\hii\\ region, the similarity to UV--excited 3.3\\mic\\ PAH emission, and through the measurements made of H$_2$ vibrational line ratios. Our data confirm this result. In particular, the strength of the v=2--1 S(1) line compared to the v=1--0 S(1) line is indicative of a non-thermal excitation method such as UV fluorescence (e.g.\\ Black \\& Dalgarno 1976, Black \\& van Dishoeck 1987). In the SW--Bar the line ratio rises above the pure fluorescent value, indicating that the density of some of the gas there is greater than critical ($\\sim 10^4$--$10^5$\\cmtwo, depending on uncertain values for the H$_2$ collisional excitation rates). In such cases collisions can re-populate the v=1 level so that the ratio of the v=1--0 to the v=2--1 S(1) lines can appear thermal (i.e.\\ $\\sim 10$, see Sternberg \\& Dalgarno 1989, Burton, Hollenbach \\& Tielens 1990). However when such ``collisional fluorescence'' is seen, density variations within the emitting region invariably give rise to significant variations in the value of this line ratio within the source (see Ryder et al.\\ 1998, Allen et al.\\ 1999). Density variations must occur in the SW--Bar of M\\,17 too. However the bulk of the gas in the SW--Bar must have density no more than critical. This can be determined from the separation between the ionization front and the excited H$_2$ there, $\\sim 5$\\,arcsec. Assuming this corresponds to an extinction of $A_v \\sim 1$ mag (i.e.\\ $\\rm N \\sim 10^{21}$ H$_2$ molecules cm $^{-2}$), it implies that the average H$_2$ number density is $\\rm n \\sim 1 \\times 10^{4} \\ cm^{-3}$ in the SW--Bar. Given the high density component that is also there, the medium must therefore be a clumpy one. In the N--Bar we will take the line ratio to be 3, as determined by Chrysostomou et at.\\ (1993). This is only a little higher than the pure fluorescent value, and suggests that the density here is about equal to the critical density. The constancy of the ratio across the N--Bar also suggests that there is also little variation in the density. This is in contrast to the clumpy SW--Bar. \\subsection{Molecular Hydrogen Formation Pumping} The most interesting result from this work is the v=6--4 O(3) emission line image of the N--bar, shown in Fig.~\\ref{fig:nlines}. This is the first time an image in such a high-excitation line of H$_2$, 31,000\\,K above ground state, has been obtained. While the line arises from the same regions as the v=1 and v=2 lines of H$_2$, its distribution within them is clearly different. Since all the lines arise from ortho (odd--J) states, the difference between them is also unlikely to result from ortho-to-para ratio variations, unless these both vary with position in the source, as well as between vibrational levels. It is also hard to see how the 6--4 O(3) line could have been mis-identified. It cannot be a line emitted from the \\hii\\ region due to the completely different morphology from the hydrogen lines. Since the line is emitted from the same regions as the lower excitation H$_2$ lines it also suggests that it is a PDR line, and not from the \\hii\\ region. Moreover, the nearby H Br10 line provided a wavelength reference point for scanning the Fabry-P\\'{e}rot etalon across the correct plate spacings for the 6--4 O(3) line. The v=6--4 O(3) line emission almost certainly arises from a region where the H$_2$ emission is dominated by UV--fluorescence. It cannot be shocked or X--ray excited since the high energy of the (v,J) = (6,1) upper level (31,000\\,K) would not be significantly populated by thermal means. The 6--4 O(3) line would be $\\sim 10^{-4}$ of the strength of the 1--0 S(1) line for thermal excitation at 2000\\,K, considerably weaker than the nearby 1--0 S(7) line, which would be about 10\\% of the 1--0 S(1) line strength. In fluorescent models, however, the v=6--4 O(3) line intensity is significant; for instance in Black \\& van Dishoeck's Model 14 it is 31\\% of the value of the 1--0 S(1) line. In the same model the 2--1 S(1) line is 56\\% the strength, and the 1--0 S(7) line 18\\%, of the 1--0 S(1) line intensity (neglecting any differential extinction between the emitting wavelengths). This is broadly consistent with the data, for which the mean fluxes of the 2--1 S(1), 1--0 S(7) and 6--4 O(3) lines are $\\sim$60\\%, 25\\% and 10\\% of the 1--0 S(1) line, respectively. However the specific predictions for the strength of the lines depend upon an additional excitation mechanism to the fluorescent cascade, formation pumping. For every H$_2$ molecule that is fluorescently excited, 15\\% lead to photodissociation (e.g.\\ Draine \\& Bertoldi 1996). In steady state PDR models, molecule destruction is balanced by molecule formation, which is believed to occur on grain surfaces (e.g.\\ Hollenbach \\& Salpeter 1971). The newly formed H$_2$ molecules are released in an excited vibrational-rotational state. However, there is little hard evidence to suggest in what state, or distribution of states, this might be. Black \\& van Dishoeck (1987) consider three possible formation models. Their `standard' Model 14 assumes that 1.5\\,eV of the binding energy is distributed in a Boltzmann distribution through the energy levels (after Black \\& Dalgarno 1976). A second model assumes that the new molecules all appear in v=14, J=0 or 1 (after Hunter \\& Watson 1978). The third formation model is based on a treatment of H$_2$ catalysis by Duley \\& Williams (1986) which predicts the molecules are ejected into the v=6 level. Naturally, there are significant differences in the predictions of the intensities for lines from v=6 between the first two and the third formation model. Depending on the formation temperature of the molecule, v=6 lines are predicted to be up to a factor $\\sim 3$ times brighter in the third than in the first formation model. This suggests that the 1--0~S(1)/6--4~O(3) ratio may fall from $\\sim 10$ to $\\sim 3$ in Black \\& van Dishoeck's (1987) Model 14, if it were to be modified so that H$_2$ formation occured in v=6, rather than into a Boltzmann distribution. We believe that formation pumping into v=6, or to a nearby level, provides the best explanation for the difference in the images between the 6--4 O(3) line and the lower excitation lines. However, it cannot simply be due to the additional component that formation pumping adds to the intensity of a line, for that would not result in a different morphology for the line. Either the rate at which formation is occurring must vary with position, or the formation spectrum itself is varying with position (for instance, due to a varying formation temperature), across the N--Bar of M\\,17\\@. In the Duley \\& Williams model, about 1.5\\,eV of the 4.5\\,eV bond energy released is assumed to remain with the H$_2$ molecule, as it is ejected from the surface of the dust grain where it formed. This puts it into an excited rotational-vibrational state, namely $\\rm v \\sim 6$. The actual formation process depends intimately on the nature of the surface of the dust grains. In their 1986 paper Duley \\& Williams considered the surfaces to be highly defected silicates, to which the molecule is moderately strongly bound. A later paper (Duley \\& Williams 1993) further considers H$_2$ formation on amorphous H$_2$O ice (where it is weakly bound) and on aromatic carbon molecules (e.g.\\ PAHs or HACs, where it is strongly bound). If the former occurs, they suggest that the H$_2$ will be released in highly-excited states (i.e.\\ the binding energy virtually all goes into internal energy in the H$_2$ molecule, so it ends up, perhaps, in v=13). In the latter case, most of the binding energy would be distributed through the many degrees of freedom of the aromatic molecule, leaving the H$_2$ in a low-excitation state. Clearly, the observations in M\\,17 support the silicate grains model for that source over the other two grain models, since the excess appears to be in v=6. The PDR models discussed above were steady state models, where the dissociation front is stationary. In them, the photodissociation rate of hydrogen molecules by far--UV radiation is balanced by their formation rate on grains. This need not be the case. The timescale for the H$_2$/H dissociation front to reach equilibrium, $\\rm t_{eq} \\sim 5 \\times 10^8/n$ years, can be long compared to variations in the radiation field (for instance, soon after a star switches on) (Hollenbach \\& Natta 1995). In M\\,17, the data constrain $\\rm t_{eq}$ to be less than $\\sim 10^4$ years. The exposure of previously shielded H$_2$ molecules to the far--UV radiation field, as the dissociation front moves further into a molecular cloud, increases the column of fluoresced gas. Thus, the H$_2$ line intensities rise from their steady state values. The 1--0 S(1) line intensity can be elevated by an order of magnitude (Hollenbach \\& Natta 1995). Moreover, since this emission is dominated by pure fluorescence (collisions can only play a minor role in redistributing the level populations), the 1--0/2--1 S(1) ratio will be $\\sim 2$. Any thermal contribution to the 1--0 S(1) line has been minimised. During this time the rate of molecule destruction must exceed that of formation. Thus the relative proportion of the formation component to the fluorescent H$_2$ line intensities will be reduced. We now consider whether this explanation can be applied towards explaining the 6--4 O(3) line emission from the N--Bar of M\\,17. In it, the line is brighter in the NE of the two emission filaments, the one furthest away (at least in projected distance) from the ionizing stars in the \\hii\\ region. Our hypothesis is that a formation component contributes to the 6--4 O(3) line intensity. This would imply that, in the weaker of the two filaments, a smaller proportion of the emission has been produced by formation pumping, compared to that produced by pure fluorescence, than in the brighter of the filaments. Thus, this suggests that the SW of the two filaments contains a non-steady state photodissociation region, where the dissociation front is moving rapidly into the molecular cloud. We can apply this interpretation to provide a rough estimate of the fraction of the 6--4 O(3) line intensity that derives directly from H$_2$ formation in a steady state PDR\\@. We assume that when the 1--0~S(1)/6--4~O(3) ratio is at it largest (i.e.\\ 23), then the UV cascade dominates its excitation. We also assume that when it is at its smallest value (i.e.\\ 5) this represents the steady state PDR\\@. This then yields a formation component that is $\\sim 80$\\% of the total 6--4 O(3) line intensity in the steady state, the remaining $\\sim 20$\\% coming from the fluorescent cascade. This is obviously a crude estimate. Without further data on other high excitation H$_2$ lines we cannot say more about the formation spectrum of H$_2$." + }, + "0112/astro-ph0112342_arXiv.txt": { + "abstract": "{ We present the results of sub-mm, mm (850$\\mu$m, 450 $\\mu$m and 1250$\\mu$m) and radio (1.4 and 4.8 GHz) continuum observations of a sample of 27 $K$-selected Extremely Red Objects, or EROs, (14 of which form a complete sample with $K<20$ and $I-K>5$) aimed at detecting dusty starbursts, deriving the fraction of UltraLuminous Infrared Galaxies (ULIGs) in ERO samples, and constraining their redshifts using the radio-FIR correlation. One ERO was tentatively detected at 1250$\\mu$m and two were detected at 1.4 GHz, one of which has a less secure identification as an ERO counterpart. Limits on their redshifts and their star forming properties are derived and discussed. We stacked the observations of the undetected objects at 850$\\mu$m, 1250$\\mu$m and 4.8 GHz in order to search for possible statistical emission from the ERO population as a whole, but no significant detections were derived either for the whole sample or as a function of the average NIR colours. These results strongly suggest that the dominant population of EROs with $K<20$ is not comprised of ULIGs like HR 10, but is probably made of radio-quiet ellipticals and weaker starburst galaxies with L$\\,<\\,$10$^{12}$ \\lsun and SFR$\\,<\\,$100 \\msunyr. ", + "introduction": "The existence of a population of extragalactic objects with extremely red infrared-optical colours has been known for a number of years (for a recent review, see Cimatti 2000\\nocite{cimdf00}). These objects were initially discovered mainly in near-infrared surveys of blind fields and quasar fields and are very faint or invisible in the optical bands (Elston et al. 1988\\nocite{err88}; McCarthy et al. 1992\\nocite{mcpw92}; Hu \\& Ridgway 1994\\nocite{hr94}). These extremely red objects (EROs) are defined as those which have $R-K$ colours $\\geq$ 5 and these tend to have $K$ magnitudes fainter than $\\sim$18. Since the time of their discovery, the nature of this population has remained a puzzle. Based on their NIR and optical photometric data, two broad classes of models are consistent with the observed red colours : (a) high redshift starbursts, red because of severe dust extinction. From the required extinction and simple dust models, these galaxies could be normal starburst galaxies or even high-$z$ counterparts of local Ultra Luminous Infrared Galaxies (ULIGs) and (b) old passively evolving ellipticals at redshifts greater than about one. The red colours of the EROs would then be explainable by a large $K$-correction and an absence of ongoing star formation. If the EROs belong to the former class, then they would be dominant sites of star formation and would be important in determining the star formation history of the universe (Cimatti et al. 1998a\\nocite{nat98}). On the other hand, if they belong to the latter class, then the volume density of these objects as a function of redshift would pose strong constraints on the models for the formation of elliptical galaxies, which range from monolithic collapse to dark matter dominated hierarchical structure formation scenarios (Daddi et al. 2000b\\nocite{dcr00} and references therein). Over the last few years, it has been possible to study in detail a handful of EROs which are bright enough to yield reliable spectra and have multi-wavelength continuum data. These studies were able to determine the nature of these EROs and there are now examples known for both starbursts (Cimatti et al. 1998a\\nocite{nat98}; Cimatti et al. 1999\\nocite{cim99}; Smail et al. 1999a\\nocite{smail99a}; Smail et al. 1999b\\nocite{smail99b}; Gear et al. 2001\\nocite{gear}) as well as for old ellipticals (Spinrad et al. 1997\\nocite{sp97}; Cimatti et al. 1999\\nocite{cim99}; Soifer et al. 1999\\nocite{soi99}; Liu et al. 2000\\nocite{liu00}) among the ERO population. Studies indicate though, that not more than $\\sim$30 \\% of EROs are starbursts (see Sect. 5 for details). Recently, independent wide-field surveys for EROs have been conducted (Daddi et al. 2000a\\nocite{daddi00}; McCarthy et al. 2000\\nocite{mccar00}; Thompson et al. 1999\\nocite{thomp99}) which have shown that these objects are strongly clustered in the sky (see also Chapman et al. 2000\\nocite{chapman00} and Yan et al. 2000\\nocite{yan00}). The surface density of these EROs after correcting for clustering, assuming that these are passively evolving ellipticals, has been shown to be consistent with pure luminosity evolution with a formation redshift greater than 2.5 (Daddi et al. 2000b\\nocite{dcr00}). HR 10 is one of the reddest EROs and is quite bright in the NIR ($I-K$=6, $K$=18.42, Graham and Dey 1996\\nocite{gd96}). \\cite{nat98} detected 850 $\\mu$m and 1250 $\\mu$m emission from this object (see also Dey et al. 1999\\nocite{dey99}). They derived a star formation rate of several hundred \\msunyr and an FIR luminosity in excess of of 10$^{12}$ \\lsunn, thus showing that HR 10 is an ULIG at its redshift of 1.44 (Graham and Dey 1996\\nocite{gd96}). At the time of the observations presented in this paper, HR 10 was the only ERO with detected submm emission and also the only ERO with a known redshift. Hence it was thought possible that a majority of EROs would be similar to HR 10 and would have observable submm and mm emission (Cimatti et al. 1998b\\nocite{cim98b}, Andreani et al. 1999\\nocite{and99}). Therefore we began a search for mm and submm continuum emission from other EROs with an aim to detect extreme starburst galaxies. Further, in order to constrain the redshifts of these objects and also to understand their nature, we also decided to search for radio continuum emission from these EROs. The radio and the FIR luminosities of nearby galaxies ($z<\\,$0.4) which are dominated by star formation are known to be highly correlated (see Condon 1992\\nocite{condon} and references therein), whereas E and S0 galaxies are known to be more radio bright than a star-forming galaxy of similar FIR luminosity (Walsh et al. 1989\\nocite{wkwk89}). The radio (mainly synchrotron and some amount of free-free) emission and the FIR (due to dust) emission have different spectral indices. Hence, assuming the local radio-FIR correlation holds at high redshift, the observed ratio of radio to FIR emission strengths can be used to determine the redshift of a star-forming galaxy (the redshift determination method and the associated error estimation are developed in Carilli and Yun (1999, 2000)\\nocite{cy99,cy00} and Blain 1999\\nocite{blain99}). So, if on the one hand, the ERO population consists primarily of starbursts, this method could be used to determine the nature of EROs and also to constrain the redshifts of these objects. If, on the other hand, the EROs are ellipticals, their IR colours would be used to constrain their properties. Since most EROs are too faint to obtain redshifts even with 10m-class telescopes or even to obtain accurate photometry over the entire optical-IR range, such complementary diagnostics become important in understanding these objects. We first describe the ERO sample, the observations and their results in Sects. 2 and 3. In Sect. 4, we derive statistical properties of the sample and estimate the fraction of starbursts and ellipticals with radio emission in Sect. 5. Those EROs with radio or mm detections are discussed in detail and their properties are derived in Sect. 6. The cosmology adopted in this paper is a flat universe with H$_{\\rm 0}$=70 \\kms~Mpc$^{-1}$ and all results are calculated for both $\\Omega_\\Lambda$=0.7 and $\\Omega_\\Lambda$=0.0. The spectral index is assumed to be $-$0.7 throughout this paper for calculating the $K$-correction for the radio emission. ", + "conclusions": "Motivated by the discovery of HR 10 and its star formation properties, a sample of EROs was observed in order to detect radio, mm and sub-mm continuum emission and constrain the redshifts and star formation rates of ULIGs in the sample. One ERO was detected at 1.4 GHz and a possible radio counterpart was identified for another ERO at the same frequency. A third was tentatively identified at 1250 $\\mu$m. Their redshifts and star forming properties were constrained using their radio-FIR spectral index but their nature could not be unambigously determined. Since the sources are faint, standard techniques to classify the sources in the ERO sample individually as ellipticals or starburst are inadequate. Weighted average flux densities were computed for the sample using measurements at the IR positions of the EROs and these values are 1.0 $\\pm$ 0.4 mJy at 850 $\\mu$m, 0.18 $\\pm$ 0.12 mJy at 1250 $\\mu$m and 3.3 $\\pm$ 4.6 $\\mu$Jy at 4.8 GHz. We find no difference within errors in the weighted average values between EROs with NIR colours above and below an assumed $I-K$ colour cut-off. If the sample is divided into ellipticals and starbursts based on the NIR two-colour diagnostic diagram, no differences in their average flux densities are seen between these groups. From the lack of detection of sub-mm emission from any of the EROs in our sample, it is now clear that dusty strong starbursts, or high redshift ULIGs, are not the dominant component of this population. If it is assumed that such an ULIG population would resemble HR 10 in their SED properties, then such galaxies cannot constitute more than about 35 \\% of the EROs. From the observed source counts of SCUBA sources and the ERO surface density, we suggest that EROs contribute negligibly to both high redshift star formation as well as to the FIRB. On the other hand, if all EROs have similar star formation properties, the average dust mass is calculated to be less than 2 $\\times$ 10$^9$ \\msun (for $\\Omega_\\Lambda$=0.7 and T$_{\\rm dust}$=20 K) and the average FIR luminosity, less than 4 $\\times$ 10$^{12}$ \\lsunn. If not more than a third of the EROs in the sample are assumed to be starbursts and the rest are assumed to be ellipticals at a median redshift of 1.5, we calculate the detection rate of ellipticals for a 0.1 mJy cut-off at 1.4 GHz to be less than 22 $\\pm$ 16 \\%. The corresponding number estimated in the local universe is $\\leq$3 \\%. Therefore it is clear that the dominant population of EROs are not ULIGs similar to HR 10 but are probably old ellipticals and weaker starbursts. The determination of the properties and redshifts of the elliptical galaxy component of the ERO population is extremely important for constraining structure formation models (Daddi et al. 2000b\\nocite{dcr00}). Towards this end, the upper limits to the radio and the sub-mm flux densities derived in this study should be used to plan future observations to detect these objects at these wavebands." + }, + "0112/astro-ph0112032_arXiv.txt": { + "abstract": "We have solved numerically the general relativistic induction equations in the interior background spacetime of a slowly rotating magnetized neutron star. The analytic form of these equations was discussed in a recent paper (Rezzolla {\\it et al} 2001a), where corrections due both to the spacetime curvature and to the dragging of reference frames were shown to be present. Through a number of calculations we have investigated the evolution of the magnetic field with different rates of stellar rotation, different inclination angles between the magnetic moment and the rotation axis, as well as different values of the electrical conductivity. All of these calculations have been performed for a constant temperature relativistic polytropic star and make use of a consistent solution of the initial value problem which avoids the use of artificial analytic functions. Our results show that there exist general relativistic effects introduced by the rotation of the spacetime which tend to {\\it decrease} the decay rate of the magnetic field. The rotation-induced corrections are however generally hidden by the high electrical conductivity of the neutron star matter and when realistic values for the electrical conductivity are considered, these corrections become negligible even for the fastest known pulsar. ", + "introduction": "\\label{intro} Irrespective of the origin of magnetic fields in neutron stars, whether produced by thermoelectric effects active in a thin layer below the star surface when the temperature is much above $10^6 \\rm{K}$ (see, for instance, Wiebicke \\& Geppert, 1996), or by a dynamo action during the earliest stages of the convective motions (see Thompson \\& Duncan, 1993), or by post core-collapse accretion of fall-back material after a supernova explosion giving rise to a neutron star, a secular decay of the magnetic field is expected as a result of the finite electrical conductivity of the stellar matter. The theoretical research in this area is intense, pushed on by the observational evidence that magnetic fields in neutron stars are decreasing with increasing spin-down age. There is now a general consensus about the possibility of improving the present knowledge of the internal structure of neutron stars by using the constraints from observations of the magnetic field decay. This justifies the effort of taking into account all of the possible factors which are supposed to play a role during the decay of the magnetic field. Particularly interesting within this context are the general relativistic corrections induced by the presence of a strongly curved background spacetime. These corrections have been investigated by a number of authors (Ginzburg \\& Ozernoy 1964, Anderson \\& Cohen 1970, Petterson 1974, Gupta {\\it et al} 1998, Konno \\& Kojima 2000) and with a number of different approaches some of which are more rigorous (Geppert {\\it et al}, 2000) than others (Sengupta 1995, 1997). In recent related works, Rezzolla {\\it et al} (2001a, 2001b) have performed a detailed analysis of Maxwell's equations in the external and internal background spacetime of a rotating magnetized conductor. As a result of this analysis, it was possible to show that in the case of finite electrical conductivity, general relativistic corrections due both to the spacetime curvature and to the dragging of reference frames are present in the induction equations. Moreover, when the stellar rotation is taken into account, each component of the magnetic field is governed by its own evolutionary law, thus removing the degeneracy encountered in the case of nonrotating spacetimes. The purpose of this paper, which is the natural extension of the work in Rezzolla {\\it et al} (2001a, hereafter paper I), is to quantify the general relativistic effects related to rotation on the evolution of the magnetic field. We have therefore solved numerically the general relativistic induction equations derived in Paper I for a relativistic polytropic star with different values of the rotation period and of the electrical conductivity. Each of the several calculations performed here benefits from the consistent solution of the initial value problem for a magnetic field which is initially permeating a perfectly conducting relativistic star. This approach avoids the use of artificial initial data and provides a more accurate solution of the induction equations. Overall, our results show that the rotation of the star and of the background spacetime introduce a {\\it decrease} in the decay rate of the magnetic field. In general, however, the rotation-induced corrections are hidden by the high electrical conductivity of the neutron star matter and are effectively negligible even for the fastest known pulsar. Also in the absence of rotation, the spacetime curvature introduces modifications to the evolution of the magnetic field when compared with the corresponding evolution in a flat spacetime. These modifications depend sensitively on both the metric functions of the interior spacetime and on the radial profile of the electrical conductivity. In the case the star is modeled as a polytrope and the electrical conductivity is assumed to be uniform in space and time, the spacetime curvature generally increases the decay rate of the magnetic field as compared to the flat spacetime case, with this increase being dependent on the compactness of the star. The paper is organized as follows: in Section 2 we discuss our treatment of the internal structure of the star in the limit of slow rotation. Section 3 is devoted to the solution of the induction equations derived in paper I, with some emphasis on the numerical aspects and in particular on the initial value problem. We show our results in Section 4, whereas Section 5 contains the conclusions. Throughout, we use a space-like signature $(-,+,+,+)$ and a system of units in which $G = c = M_{\\odot} = 1$ (However, for those expressions of astrophysical interest, we have written the speed of light explicitly.). Partial spatial derivatives are denoted with a comma. ", + "conclusions": "\\label{conc} In a recent paper, Rezzolla {\\it et al}~(2001a) have considered the general relativistic description of the electromagnetic fields of a slowly rotating, magnetized and misaligned neutron star. If the stellar medium has a finite electrical conductivity it was shown that the stellar rotation removes the degeneracy in the evolution equations for the magnetic field and that three distinct induction equations need to be solved to account for the decay of the stellar magnetic field. In this paper we have solved numerically the general relativistic induction equations derived in paper I, investigating the effects of different rotation rates, different inclination angles between the magnetic moment and the rotation axis, as well as different values of the electrical conductivity. The aim of these numerical calculations is that of quantifying the corrections induced by general relativistic effects (both due to spacetime curvature and to the stellar rotation) on the evolution of the magnetic field of a slowly rotating neutron star. In order to single out purely general relativistic effects from those due to the microphysics of the Ohmic dissipation, we have considered a simplified physical description of the neutron star. In particular, the star has been modelled as a polytrope rotating with a fiducial period of one millisecond, the electrical conductivity has been considered to be uniform inside the star and we have not included a treatment to consider the evolution of the stellar rotation and temperature (see Page {\\it et al} 2000). On the other hand, special attention has been paid to a consistent solution of the initial value problem and we have considered as initial magnetic field the stationary solution of the general relativistic Maxwell equations. In this way we have avoided the use of initial magnetic field configurations that are only approximate solutions of the Maxwell equations (i.e. solutions of the Maxwell equations only in the limit of flat spacetime). Besides eliminating an initial error during the initial stages of the magnetic field decay, our prescription for the initial value problem also provides a more accurate solution of the Maxwell equations. The results of our computations have shown that there exist general relativistic, rotation-induced corrections to the evolution of the magnetic field. These effects generally produce a {\\it decrease} in the rate of magnetic field decay. However, their contribution is masked by the high value of the electrical conductivity in realistic neutron stars and can be neglected in general. Our calculations also indicate that general relativistic effects not induced by the stellar rotation can modify the time evolution of the magnetic field in a magnetized star. Such effects are closely related to the properties of the spacetime internal to the star and for a polytropic stellar model with uniform electrical conductivity these effects generally {\\it increase} the decay rate of the field. The validity of this conclusion is however limited. Density gradients are in fact expected in a realistic star and these will affect the behaviour of the electrical conductivity which, in turn, will influence the decay of the magnetic field. Our conclusions are that the general relativistic evolution of the magnetic field in rotating neutron stars can be studied with confidence already in a nonrotating background spacetime. However, the role of a curved background spacetime on the decay of the magnetic field can be fully assessed only when the details of both a realistic equation of state and of a realistic electrical conductivity are carefully taken into account. This will be the subject of future work." + }, + "0112/astro-ph0112518_arXiv.txt": { + "abstract": "At least three pulsars in supernova remnants were detected at E $>$ 100 MeV by EGRET on the Compton Gamma Ray Observatory. Efforts to search for additional pulsars in the EGRET data have been unsuccessful due to limited statistics. An example is the recently-discovered radio pulsar J2229+6114, where efforts to search the EGRET data using several different methods failed to find significant evidence of pulsation. The GLAST Large Area Telescope (LAT) will have a much greater effective area and a narrower point-spread function than EGRET. In addition, the field of view will be more than 4 times larger than EGRET's, and the LAT will scan to avoid occultation by the earth, increasing by a large factor the total number of photons detected. The greater rates of photons from pulsar candidates and better discrimination of diffuse interstellar emission will enhance the sensitivity of pulsation searches. These improvements also offer the prospect of resolving point sources from extended emission in some SNR to define the nature of the associations of EGRET sources with SNR. Further, work with the GLAST LAT will benefit from ongoing multiwavelength studies (e.g., for RX 1836.2+5925) that provide specific candidate targets for gamma-ray studies. ", + "introduction": "At least three of the rotation-powered pulsars seen in high-energy gamma rays by EGRET (Crab, Vela, PSR B1951+32) are neutron stars in SNR. A fourth, PSR B1706$-$44, would be expected from its age to have a visible SNR, but the observational evidence is unclear. (Dodson, Gvamaradze, these proceedings.) Because EGRET and previous gamma-ray telescopes lacked the resolution to distinguish neutron star emission from nebular emission, the only certain detections are based on pulsed emission. \\begin{figure} \\centerline{\\psfig{figure=thompsond_1.eps,height=4.5cm,bbllx=15pt,bblly=15pt,bburx=570pt,bbury=370pt,clip=.}} \\caption{Possible EGRET light curve for PSR J2229+6114 in Viewing Period 34, based on a search of periods and period derivatives consistent with the later radio measurements.} \\end{figure} ", + "conclusions": "" + }, + "0112/astro-ph0112204_arXiv.txt": { + "abstract": "We investigate the relationship between \\HI, H$_2$, and the star formation rate (SFR) using azimuthally averaged data for seven CO-bright spiral galaxies. Contrary to some earlier studies based on global fluxes, we find that \\sigsfr\\ exhibits a much stronger correlation with \\sightwo\\ than with \\sighi, as \\sighi\\ saturates at a value of $\\sim$ 10 \\Msol\\ pc$^{-2}$ or even declines for large \\sigsfr. Hence the good correlation between \\sigsfr\\ and the total (\\HI+H$_2$) gas surface density \\siggas\\ is driven by the molecular component in these galaxies, with the observed relation between \\sigsfr\\ and \\sightwo\\ following a Schmidt-type law of index $n_{\\rm mol}$$\\approx$0.8 if a uniform extinction correction is applied or $n_{\\rm mol}$$\\approx$1.4 for a radially varying correction dependent on gas density. The corresponding Schmidt law indices for \\siggas\\ vs.\\ \\sigsfr\\ are 1.1 and 1.7 for the two extinction models, in rough agreement with previous studies of the disk-averaged star formation law. An alternative to the Schmidt law, in which the gas depletion timescale is proportional to the orbital timescale, is also consistent with the data if radially varying extinction corrections are applied. We find no clear evidence for a link between the gravitational instability parameter for the gas disk ($Q_g$) and the SFR, and suggest that $Q_g$ be considered a measure of the gas fraction. This implies that for a state of marginal gravitational stability to exist in galaxies with low gas fractions, it must be enforced by the stellar component. In regions where we have both \\HI\\ and CO measurements, the ratio of \\HI\\ to H$_2$ surface density scales with radius as roughly $R^{1.5}$, and we suggest that the balance between \\HI\\ and H$_2$ is determined primarily by the midplane interstellar pressure. These results favor a ``law'' of star formation in quiescent disks in which the ambient pressure and metallicity control the formation of molecular clouds from \\HI, with star formation then occurring at a roughly constant rate per unit H$_2$ mass. ", + "introduction": "Understanding the factors which control the star formation rate (SFR) is of fundamental importance to studies of the interstellar medium (ISM) and galaxy evolution. The most important factor, of course, is the availability of cold neutral gas, especially molecular gas. Over the past several decades, observations of the CO(1--0) line have established a clear link between star formation in the Galaxy and molecular gas \\citep[e.g.,][]{Loren:73}, and even in environments where star formation is observed without CO emission, such as dwarf irregular galaxies, a plausible case can be made that molecular clouds are weak CO emitters due to low metal abundances \\citep{Wilson:95,Taylor:98}. On the other hand, \\citet{KC:89} has shown that the disk-averaged \\sigsfr\\ is much better correlated with \\sighi\\ than with \\sightwo\\ inferred from CO observations, a surprising result given the very different radial profiles of \\HI\\ and H$\\alpha$ emission in galaxies. Similar conclusions have been reached by other studies based on global averages \\citep{Deharveng:94,Boselli:94,Casoli:96}. Thus, it is not obvious whether the \\HI\\ or H$_2$ (as traced by CO) content will be a better predictor of the SFR on large scales. \\citet{KC:89} argues that it is the {\\it total} (\\HI+H$_2$) gas surface density which shows the best correlation with the SFR per unit area. A quantitative prescription for the SFR has been provided by \\citeauthor{KC:98a} (1998a, hereafter K98), who finds that the total gas surface density, averaged over the optical disk, is related to the SFR surface density by a \\citet{Schmidt:59} law, \\begin{equation} \\sigsfr \\propto (\\siggas)^n\\;. \\label{eqn:schmidt} \\end{equation} where $n \\approx 1.4$. Kennicutt's formulation of the star formation law has been widely applied in ``semi-empirical'' models of galaxy evolution \\citep[e.g.,][]{Boissier:99,Tan:99,vdB:00}. While the Schmidt law appears to be valid at high gas surface densities, \\citet{KC:89} and \\citet{Martin:01} have presented evidence for a threshold density for star formation based on the gravitational instability parameter $Q$. Below the threshold density, which observationally is of order 5--10 \\Msol\\ pc$^{-2}$, the Schmidt law breaks down and star formation is strongly suppressed. The general applicability of a $Q$ threshold, however, remains controversial \\citep{Thornley:95,Hunter:98,Ferguson:98b}. \\begin{table*}[b] \\begin{center} \\caption{Properties of the Sample Galaxies\\label{tbl:props}} \\bigskip \\begin{tabular}{cccccccc} \\hline\\hline Name & RA (2000) & DEC & Morph. & $B_T^0$\\tablenotemark{(a)} & $R_{25}$\\tablenotemark{(b)} & $V_{\\odot}$\\tablenotemark{(c)} & Dist. \\\\ & hh:mm:ss & dd:mm:ss & Type & (mag) & (arcsec) & (km/s) & (Mpc) \\\\[0.5ex] \\hline\\hline NGC 4321 & 12:22:54.9 & 15:49:21 & SAB(s)bc & 9.98 & 220 & 1571 & 16\\tablenotemark{(d)}\\\\ NGC 4414 & 12:26:27.1 & 31:13:24 & SA(rs)c? & 10.62 & 110 & 716 & 19\\tablenotemark{(e)}\\\\ NGC 4501 & 12:31:59.2 & 14:25:14 & SA(rs)b & 9.86 & 210 & 2281 & 16\\tablenotemark{(d)}\\\\ NGC 4736 & 12:50:53.1 & 41:07:14 & (R)SA(r)ab & 8.75 & 340 & 308 & 4.2\\tablenotemark{(f)}\\\\ NGC 5033 & 13:13:27.5 & 36:35:38 & SA(s)c & 10.21 & 320 & 875 & 17.5\\tablenotemark{(g)}\\\\ NGC 5055 & 13:15:49.3 & 42:01:45 & SA(rs)bc & 9.03 & 380 & 504 & 10\\tablenotemark{(g)}\\\\ NGC 5457 & 14:03:12.5 & 54:20:55 & SAB(rs)cd & 8.21 & 870 & 241 & 7.4\\tablenotemark{(h)}\\\\ \\hline\\hline \\end{tabular} \\end{center} \\tablenotetext{a}{Total blue magnitude, corrected for Galactic and internal absorption and redshift, from RC3 \\citep{RC3}.} \\tablenotetext{b}{Semimajor axis at 25.0 mag arcsec$^{-2}$, from RC3.} \\tablenotetext{c}{Heliocentric velocity, from NED.} \\tablenotetext{d}{Cepheid distance to NGC 4321 from \\citet{Ferrarese:96} adopted for Virgo cluster.} \\tablenotetext{e}{Cepheid distance from \\citet{TurnerA:98}.} \\tablenotetext{f}{Redshift distance using $H_0$ = 75 \\kms\\ Mpc$^{-1}$.} \\tablenotetext{g}{Distance from Mark III data \\citep{Willick:97} using $H_0$=75.} \\tablenotetext{h}{Cepheid distance from \\citet{Kelson:96}.} \\end{table*} As noted by K98, an alternative formulation of the star formation law, where the star formation timescale is proportional to the orbital timescale, \\begin{equation} \\sigsfr \\propto \\siggas \\Omega\\;, \\end{equation} is also consistent with the disk-averaged data. Such a law is predicted by models in which spiral arms play an important role in triggering star formation \\citep{Wyse:89}, or in which star formation is self-regulated to yield a constant value of $Q$ \\citep{Larson:88,Silk:97}. Distinguishing between this law and a Schmidt law requires {\\it spatially resolved} measurements of gas densities and orbital timescales (i.e., rotation curves). A third prescription for the SFR, arising from CO studies by J. Young and collaborators \\citep{Devereux:91,Young:96,Rownd:99}, states that the SFR is roughly proportional to the mass of molecular gas, so that their ratio (commonly referred to as the {\\it star formation efficiency}, SFE) is roughly constant. In particular, \\citet{Rownd:99} find that the SFE within the disks of spiral galaxies exhibits no strong radial gradients, under the assumption of a constant CO-to-H$_2$ conversion factor. An SFE that is independent of galactocentric radius is also consistent with studies of molecular clouds in our Galaxy \\citep{Mead:90,Wouter:88}. If star formation in disk galaxies can be generically decomposed into two processes, the formation of giant molecular clouds (GMCs) from \\HI\\ clouds and the formation of stars within GMCs \\citep[e.g.,][]{Tacconi:86}, then the latter process occurs at a rate that is largely independent of location within a galaxy, and may well be universal in all disk galaxies. Previous studies of the star formation law in external galaxies have tended to concentrate on large samples observed at the low resolution provided by single-dish radio telescopes. Although such samples are valuable for characterizing the global properties of galaxies, detailed study of individual objects at high spatial resolution is essential for understanding the physical basis of the star formation law and reconciling the various empirical descriptions. In this paper we undertake such an investigation, comparing recent CO data from the BIMA Survey of Nearby Galaxies (BIMA SONG) with previously published \\HI\\ and H$\\alpha$ data. Our high-resolution CO maps include single-dish data, which is essential when accurate flux measurements are required for extended sources. We employ azimuthally averaged radial profiles in order to improve the signal-to-noise ratio and average over time-dependent effects. The organization of this paper is as follows. In \\S\\ref{obs} we describe our galaxy sample, observing parameters, and data reduction procedures. \\S\\ref{analysis} describes how we produced radial profiles from the intensity images and our corrections for H$\\alpha$ extinction. \\S\\ref{schmidt} presents our results for the relation between \\sigsfr\\ and \\siggas, as well as the relations between \\sigsfr\\ and \\sighi\\ and \\sightwo\\ individually. In \\S\\ref{crit} we compare \\siggas\\ with the threshold surface density discussed by \\citet{KC:89}, and in \\S\\ref{hitoh2} we investigate the relationship between the radial profiles of atomic and molecular gas. In \\S\\ref{disc} we show how our results can be interpreted in terms of a nearly constant star formation efficiency for molecular gas, with molecular cloud formation in turn linked to the average ISM pressure. Finally, in \\S\\ref{conc} we summarize our results. In subsequent papers of this series, we will analyze the CO and \\HI\\ kinematics of the galaxies in our sample for evidence of radial inflows (Wong, Blitz, \\& Bosma 2002, Paper II) and evaluate whether the observed gas depletion times and oxygen abundances are consistent with closed-box evolution (Wong \\& Blitz 2002, Paper III). ", + "conclusions": "\\label{conc} We have investigated the relationship between gas content and star formation rate, commonly termed the star formation law, in seven nearby galaxies using CO, HI, and H$\\alpha$ images at resolutions of $\\sim$20\\arcsec\\ or higher. An important advance is the recent availability of CO data that combines the resolution of an interferometer with the sensitivity to large-scale structure afforded by a single-dish telescope. Although our sample is biased toward luminous, molecule-rich galaxies, none of the galaxies display signs of strong interactions. We find that the correlation of the azimuthally averaged SFR surface density \\sigsfr\\ with \\sighi\\ is much poorer than with \\sightwo, contrary to the results of studies based on integrated fluxes. Whereas a roughly linear relation exists between \\sigsfr\\ and \\sightwo, consistent with a constant star formation efficiency for the molecular gas, \\sighi\\ tends to reach a maximum value of $\\sim$10 \\Msol\\ pc$^{-2}$, or even decline in regions of high SFR. Thus the star-forming gas in these galaxies exists predominantly in molecular form, although the situation may be different in low-mass galaxies. Despite the poor correlation of \\sigsfr\\ with \\sighi, the dominance of H$_2$ in these galaxies leads to a strong correlation between \\sigsfr\\ and total gas surface density \\siggas\\ when azimuthally averaged, in rough agreement with the global Schmidt law derived by K98. The index of the power law is given by $n \\approx 1.1$ if the extinction at H$\\alpha$ is assumed to be uniform with radius, but steepens to $n \\approx 1.7$ if the mean extinction is assumed to vary according to the gas surface density. A star formation law of the form $\\sigsfr \\propto \\siggas\\Omega$ is also consistent with the data, but only if the $N_{\\rm H}$-dependent extinction corrections are applied. We suggest that the observed Schmidt law results from an effective dependence of the molecular fraction on \\siggas\\ by way of the interstellar pressure, coupled with a roughly linear relation between \\sigsfr\\ and \\sightwo. The gravitational instability parameter $Q_g$ is found to be $\\sim 1$ in several galaxies but is clearly $>1$ in the two galaxies with the smallest gas fractions, NGC 4736 and 5457. We suggest that $Q_g$ is best considered a measure of the gas fraction, and that when discussing a range of galaxy types, the contribution of the stellar component to the disk instability should not be neglected. While a combined (gaseous and stellar) instability may ultimately provide a better description of star formation thresholds, alternative explanations that attribute such thresholds to the inability to form cold, dense clouds remain attractive as well. Finally, over the limited range in radius where our CO and \\HI\\ data overlap, we find that the relation $\\sighi/\\sightwo \\propto R^{1.5}$ holds remarkably well within galaxies, although the proportionality constant varies among galaxies. We show that such a relation supports the view that the interstellar pressure plays the dominant role in determining the \\HI-H$_2$ balance. These results have demonstrated the importance of including the molecular component in studying the star formation law. Moreover, considering just the total gas density rather than the \\HI\\ and H$_2$ surface densities separately may obscure underlying physical processes that are essential to star formation. In a future study, we will develop these ideas using the much larger CO database provided by the BIMA SONG project, supplemented by archival and recently obtained \\HI\\ synthesis imaging. We also emphasize that significant advances will be made possible by improved measurements of extinction within galaxies to enable more accurate estimates of the SFR, and by studies of the vertical structure of edge-on galaxies, which will provide further insight into the \\HI-H$_2$ transition and its relationship to the ISM pressure." + }, + "0112/astro-ph0112296_arXiv.txt": { + "abstract": "We use {\\em Far Ultraviolet Spectrocopic Explorer} (FUSE) observations to study interstellar absorption along the line of sight to the white dwarf WD~1634-573 ($d=37.1\\pm 2.6$~pc). Combining our measurement of D~I with a measurement of H~I from {\\em Extreme Ultraviolet Explorer} data, we find a D/H ratio toward WD~1634-573 of ${\\rm D/H}=(1.6\\pm 0.5)\\times 10^{-5}$. In contrast, multiplying our measurements of ${\\rm D~I/O~I}=0.035\\pm 0.006$ and ${\\rm D~I/N~I}=0.27\\pm 0.05$ with published mean Galactic ISM gas phase O/H and N/H ratios yields ${\\rm D/H}_{O}=(1.2\\pm 0.2)\\times 10^{-5}$ and ${\\rm D/H}_{N}=(2.0\\pm 0.4)\\times 10^{-5}$, respectively. Note that all uncertainties quoted above are 2$\\sigma$. The inconsistency between ${\\rm D/H}_{O}$ and ${\\rm D/H}_{N}$ suggests that either the O~I/H~I and/or the N~I/H~I ratio toward WD~1634-573 must be different from the previously measured average ISM O/H and N/H values. The computation of ${\\rm D/H}_{N}$ from D~I/N~I is more suspect, since the relative N and H ionization states could conceivably vary within the LISM, while the O and H ionization states will be more tightly coupled by charge exchange. ", + "introduction": "The deuterium-to-hydrogen (D/H) ratio is of central importance to many areas of astrophysics. In the standard Big Bang cosmology, the primordial D/H value places strong constraints on the total baryon content in the universe. Since the deuterium abundance is believed to decrease with time due to the destruction of deuterium in stellar interiors, measurements of the local D/H value in our galaxy provide a lower limit for the primordial D/H value, while measurements of D/H in the Lyman-$\\alpha$ forest provide an estimate that may be closer to the true primordial D/H ratio. Comparing the local and Lyman-$\\alpha$ forest values provides a measure of how much stellar processing of material has occurred during the lifetime of the Milky Way Galaxy. Measurements of D/H in different locations in the Galaxy can provide critical tests of Galactic chemical evolution models. The {\\em Copernicus} satellite provided the first accurate D/H measurements within the Galactic ISM. \\citet{jbr73} measured ${\\rm D/H}=(1.4\\pm 0.2)\\times 10^{-5}$ toward $\\beta$~Cen, and \\citet{dgy76} quoted an average value of ${\\rm D/H}=(1.8\\pm 0.4)\\times 10^{-5}$ for several lines of sight observed by {\\em Copernicus}. The Local Interstellar Cloud (LIC) within which the Sun resides appears to have ${\\rm D/H}=(1.5\\pm 0.1)\\times 10^{-5}$ based on {\\em Hubble Space Telescope} (HST) measurements along many lines of sight through the LIC \\citep{jll98}, with no evidence for any significant variation within this small cloud \\citep[$\\sim 7$ pc across based on a compilation of LIC measurements;][]{sr00}. Estimating the primordial value from the Lyman-$\\alpha$ forest has proved to be more difficult, although these measurements do suggest a higher D/H in the Lyman-$\\alpha$ forest compared with the LIC, as one would expect \\citep[e.g.,][]{jkw97,sb98,dt99,dk00}. For a more thorough review of these intergalactic D/H measurements and how they relate to Galactic D/H, see \\citet{hwm01}. However, the LIC D/H value quoted above is also problematic, as it measures only very local interstellar gas. Values of D/H elsewhere in the Galaxy could in principle be much different, since material in different parts of the Galaxy has undergone varying amounts of stellar processing. Evidence for variations of D/H within the LISM have in fact been found for longer lines of sight sampling many clouds, based on data from HST and the IMAPS instrument \\citep{avm98,ebj99,gs00,ml01}. Measuring the variation of D/H within the Galaxy is important both for determining the degree of mixing of ISM material within the Milky Way, and for determining a better estimate for the mean present day Galactic D/H value. Unfortunately, HST is not able to obtain measurements of D/H for long lines of sight through the Galaxy, because for column densities greater than $5\\times 10^{18}$ cm$^{-2}$ the H~I Ly$\\alpha$ absorption becomes so broad that it completely obscures the D I absorption. Measuring deuterium at different locations throughout our Galaxy is one of the primary missions of the {\\em Far Ultraviolet Spectroscopic Explorer} (FUSE). Unlike HST, FUSE can observe the higher lines of the H~I Lyman series, which allows deuterium to be observed for higher ISM column densities. This paper is one of a series reporting the first deuterium measurements made with FUSE \\citep{sdf01,gh01,jwk01,nl01,ml01,gs01}, including a summary paper of these initial results \\citep{hwm01}. The line of sight analyzed here is that toward the white dwarf WD~1634-573 (=HD~149499B). Table~1 lists properties of our target star, which is a very metal poor DO white dwarf unlikely to have photospheric metal lines contaminating the ISM absorption lines in which we are interested \\citep{sd99}. There is a K0~V companion star (HD~149499A) that will be in the FUSE aperture for our observations. However, even though this star is believed to be quite active \\citep{sj97}, emission from this star will not be detectable against the very bright continuum of the hot white dwarf, based on extrapolations of H~I Lyman line fluxes observed for other active single K dwarfs. With a path length of only $37.1\\pm 2.6$~pc \\citep{macp97}, the WD~1634-573 line of sight does not exceed the distance scale of deuterium measurements accessible to HST. However, the H~I column density estimated from {\\em Extreme Ultraviolet Explorer} (EUVE) data is $N_{\\rm H}=(7\\pm 2)\\times 10^{18}$ cm$^{-2}$ \\citep{rn96}, too high for deuterium to be detectable in Ly$\\alpha$, so this line of sight {\\em does} extend deuterium measurements to higher columns than are measurable with HST. ", + "conclusions": "Because there are many D~I, N~I, and O~I lines of varying strengths, we believe our column density measurements for these species are quite accurate, with 2$\\sigma$ uncertainties of only 0.05--0.07 dex (see Table~4). The uncertainty roughly doubles when we have only one or two unsaturated lines to work with, as is the case for P~II, S~II, Ar~I, and Fe~II. When we have only saturated lines to work with, the uncertainties are much higher. Thus, the H~I, C~II, C~III, and N~II column densities have uncertainties in the 0.3--0.7 dex range. This problem with saturated lines exists for H~I despite there being many lines to work with, at least one of which (H~I Ly$\\beta$) is at least partly out of the flat part of the curve of growth (see above). Nevertheless, our analysis suggests that FUSE data cannot be used to derive a precise H~I column density toward WD~1634-573. The situation will likely be better for higher column density lines of sight with at least Ly$\\beta$ and Ly$\\gamma$ completely out of the flat part of the curve of growth. However, \\citet{ml01} express concern that undetected hot H~I components, such as the heliospheric and astrospheric absorption previously observed toward nearby cool stars \\citep*{bew00}, could potentially affect any measurement of H~I column densities from the Lyman absorption lines, making it even harder to assess systematic errors for these analyses. Because our H~I measurement ($\\log {\\rm N(H~I)}=18.6\\pm 0.4$) is not very precise, we must look elsewhere for an H~I measurement that we can use to compute an accurate D/H value for this line of sight. \\citet{jbh98} estimated $\\log {\\rm N(H~I)}\\approx 18.8$ from their analysis of IUE observations of the Ly$\\alpha$ line. \\citet{rn95} found $\\log {\\rm N(H~I)}=19.0\\pm 0.4$ from ORFEUS observations of the same H~I Lyman lines that we have analyzed using FUSE. Our results are consistent with ORFEUS, but the error bars are large for both measurements. The most accurate measurement appears to be from EUVE observations of H~I Lyman continuum absorption shortward of 912~\\AA. In a simultaneous analysis of both ORFEUS and EUVE data, which allowed both photospheric parameters (e.g., $T_{eff}$ and $\\log g$) and the interstellar H~I column density to vary, \\citet{rn96} finds ${\\rm N(H~I)}=(7\\pm 2)\\times 10^{18}$ cm$^{-2}$ (i.e., $\\log {\\rm N(H~I)}=18.85\\pm 0.12$). See also \\citet{sj97} for more details on the photospheric analysis. Based on this H~I value and our D~I measurement ($\\log {\\rm N(D~I)}=14.05\\pm 0.05$), which are both quoted with 2$\\sigma$ uncertainties, we find that ${\\rm D/H}=(1.6\\pm 0.5)\\times 10^{-5}$ toward WD~1634-573. This value is consistent with the LIC value of ${\\rm D/H}=(1.5\\pm 0.1)\\times 10^{-5}$ \\citep{jll98}. The Ar~I/H~I abundance ratio, $\\log {\\rm Ar~I/H~I}=-5.59\\pm 0.17$, is close to the B star Ar abundance of $\\log {\\rm Ar/H}=-5.50$ \\citep{deh90}. This is in contrast to the results of \\citet{ebj00}, who found significantly lower Ar~I/H~I ratios toward 3 stars ($\\log {\\rm Ar~I/H~I}\\approx -5.9$), which they attributed to a higher ionization state of Ar compared with H due to the high photoionization cross section of Ar. This provides important evidence that the LISM is in fact in photoionization equilibrium. Perhaps the WD~1634-573 line of sight is more shielded from photoionization, allowing the ionization state of Ar to be closer to that of H. The apparently high N~II/N~I ratio (see Table~4) might suggest otherwise, although uncertainty in the N~II column is large. Due to charge exchange interactions, the ionization states of D, H, and O are closely coupled. This means that the total D/O gas phase abundance ratio is well approximated by the D~I/O~I ratio. Charge exchange is also important between D, H, and N, although to a lesser extent, so ${\\rm D/N}\\approx {\\rm D~I/N~I}$. Our measurements of D~I, O~I, and N~I in Table~4 yield ${\\rm D~I/O~I}=0.035\\pm 0.006$ and ${\\rm D~I/N~I}=0.27\\pm 0.05$ (2$\\sigma$ errors). These ratios can be used to estimate D/H when we multiply them by previously measured values for the gas phase O/H and N/H ratios in the ISM. \\citet{dmm98} used GHRS observations of the O~I $\\lambda$1356 lines of 13 OB stars that are more distant ($d=100-1000$ pc) than WD~1634-573 to estimate ${\\rm O/H}=(3.43\\pm 0.15)\\times 10^{-4}$, with no evidence for variation within the ISM. Note that we have increased the published O/H value by 7.5\\% to account for a revised f value for O~I $\\lambda$1356 suggested by \\citet{dew99}. Likewise, \\citet{dmm97} find ${\\rm N/H}=(7.5\\pm 0.4)\\times 10^{-5}$ toward a similar sample of 7 OB stars, also with no evidence of variation. Note that these measurements of O/H and N/H only consider the column densities of O~I, N~I, H~I, and H$_{2}$, and do not take into account the column densities of the ionized species O~II, N~II, or H~II. Since ${\\rm O~I/O~II}={\\rm H~I/H~II}$ is a good approximation, the O/H value should be relatively free from inaccuracies induced by ionization state issues, but the N/H value is potentially susceptible to differences and variations in the relative ionization states of N and H (see above). Note also that the cited uncertainties in O/H and N/H are 1$\\sigma$ standard deviations of the mean rather than simple standard deviations about the mean, meaning that the scatter of the individual O/H and N/H measurements is actually larger than suggested by the quoted uncertainties, potentially allowing for a larger variation of O/H and N/H within the ISM than is suggested by those uncertainties. In any case, multiplying our D~I/O~I measurement with the \\citet*{dmm98} mean O/H value results in ${\\rm D/H}_{O}=(1.2\\pm 0.2)\\times 10^{-5}$, while multiplying our D~I/N~I measurement with the \\citet*{dmm97} mean N/H value results in ${\\rm D/H}_{N}=(2.0\\pm 0.4)\\times 10^{-5}$ (2$\\sigma$ errors). These two D/H values do not overlap, suggesting that the either the O~I/H~I and/or N~I/H~I ratio toward WD~1634-573 is not consistent with the \\citet*{dmm97,dmm98} mean O/H and N/H values. As suggested above, the N~I/H~I ratio is more suspect due to the possibility of variations in the relative ionization states of N and H within the LISM. Furthermore, the N~I/H~I value measured toward G191-B2B is clearly lower than the \\citet*{dmm97} N/H value \\citep{avm98,ml01}. A more extensive discussion of these issues is provided by \\citet{hwm01}. Note that the ${\\rm D/H}_{O}$ measurement is a bit lower than the accepted LIC value of ${\\rm D/H}=(1.5\\pm 0.1)\\times 10^{-5}$." + }, + "0112/astro-ph0112319_arXiv.txt": { + "abstract": "{We report on the analysis of 2MASS near-infrared data of a sample of low-mass stars and brown dwarfs in the $\\sigma$\\,Orionis cluster. Youth and cluster membership have been spectroscopically confirmed using the Li~{\\sc i} spectral line. We find little evidence in the $JHK_{\\rm s}$ colour-colour diagram for near-infrared excess emission for these cluster members. By comparison with model expectations, at most 2 out of 34 stars show $(H-K_{\\rm s})$ colour consistent with a near-infrared excess. This scarcity of near-infrared signatures of circumstellar disks in the lower-mass and substellar regimes of this cluster contrasts with findings in younger clusters, hinting at an age dependence of the disk frequency. Taking into account the apparent cluster age, our result supports the idea of a relatively fast (few Myr) disk dissipation and extends this conclusion to the substellar regime. We also find some evidence that, in this cluster, the disk frequency as measured by the K$_{\\rm s}$-band excess may be mass dependent. ", + "introduction": "Star formation is believed to start with the fragmentation of rotating and magnetized clouds of interstellar gas, and further gravitational collapse forms stars across the whole mass range. Flattened disk-like structures arise during this star formation process (see review by Evans \\cite{evans99}), and circumstellar disks are common around young stars. Even though it is thought that brown dwarfs are just the substellar counterparts of stars, forming as outlined above, it has also been proposed that brown dwarfs form either through instabilities in circumstellar disks (e.g.\\ Pickett et al.\\ \\cite{pickett00}) or as stellar embryos ejected from multiple proto-stellar systems (Reipurth \\& Clarke \\cite{reipurth01}). Disk dissipation timescales, together with the environment of the young stellar population, are of particular importance, in the context of pre-main-sequence stellar evolution and of planet formation processes. Haisch et al.\\ (\\cite{haisch01}) have found evidence for an age dependence of the disk frequency from the analysis of low-mass stars in several young clusters in OB associations and star forming regions. Even though the importance of the cluster environment is not yet clear, it has been proposed that photo-evaporation by the massive OB stars plays a role in disk dissipation (e.g.\\ Johnstone et al. \\cite{johnstone98}; Bally et al.\\ \\cite{bally00}). It is also relevant to probe how the observed disk properties extend into the substellar regime. \\object{$\\sigma$\\,Orionis} is a member of the Orion\\,OB1b association, with an estimated age of 1.7$-$7\\,Myr and a distance modulus of 7.8$-$8\\,mag (e.g.\\ Warren \\& Hesser \\cite{warren78}; Brown et al.\\ \\cite{brown94}) --- the Hipparcos distance to this star is 352$^{+166}_{-85}$\\,pc. ROSAT observations and follow-up photometry (Wolk \\cite{wolk96}) revealed a cluster of low-mass young stellar objects near $\\sigma$\\,Ori. The reddening towards this massive O-type star is low ($E(B-V)$\\,=\\,0.05\\,mag, e.g.\\ Brown et al.\\ \\cite{brown94}), suggesting that the young stellar population is affected by low extinction. Kenyon et al.\\ (\\cite{kenyon01}) have confirmed both youth and cluster membership for a sample of low-mass and substellar objects within 30\\arcmin\\ of $\\sigma$\\,Ori. We analyse 2MASS near-infrared photometric data for their sample of stars. The $JHK_{\\rm s}$ colour-colour diagram of these cluster members shows little evidence for near-infrared excess. By comparing observed $(H-K_{\\rm s})$ with model expectations, we find that only 2 objects out of 34 show an infrared excess suggestive of circumstellar disks. This is in striking contrast to high disk frequencies determined for the low-mass populations in younger clusters (Trapezium, NGC\\,2024 and NGC\\,2264, Haisch et al.\\ (\\cite{haisch01}) and references therein). Our result provides valuable support for the age dependence of disk frequency in low-mass stars and may indicate that this dependence continues into the substellar regime. ", + "conclusions": "Our sample of low-mass $\\sigma$\\,Ori cluster members shows an extremely low fraction of objects with near-infrared excess: at most 2 out of 34. As near-infrared excesses in pre-main-sequence stars are normally attributed to circumstellar disks, we interpret this result as a 6\\% disk frequency in the $\\sigma$\\,Ori cluster (7\\% if the 4 apparently older stars are removed from the sample as possible contaminants). Using an X-ray selected sample ($M \\gsim$ 0.5\\,M$_{\\odot}$), Wolk (\\cite{wolk96}) proposed a 12\\% (6/49) disk frequency, using H$\\alpha$ emission as an indirect indicator of circumstellar material (emission equivalent width $>$\\,10\\,\\AA). X-ray selection might bias the sample against stars with disks (Stelzer \\& Neuh\\\"{a}user \\cite{stelzer01}). However, H$\\alpha$ samples are biased towards stars with high accretion rates. Stassun et al.\\ (\\cite{stassun99}) show that some low-mass stars (in Orion) with K-band excess do not show strong H$\\alpha$ emission. For both these reasons, Wolk's estimated disk fraction must be a lower limit to the fraction one would estimate from near-infrared photometry. Therefore, there is some evidence, tempered by small number statistics, that the disk frequency (as measured by K-band excess) decreases towards lower masses in the $\\sigma$\\,Ori association. As the near-infrared colours trace the warmer dust, the magnitude of the near-infrared excess depends on the position and temperature of the inner disk boundary. Towards very low masses, the stellar effective temperature is critical for the disk detection efficiency at these wavelengths. Natta \\& Testi (\\cite{natta01}) show that classical T Tauri disk models describe well the observed spectral energy distribution (SED) for brown dwarf candidates. They suggest that, in the substellar regime, the inner parts of such disks might not be hot enough to produce detectable near-infrared excess. Indeed, Comer\\'{o}n et al.\\ (\\cite{comeron00}) find that several low-mass stars and brown dwarfs (0.03\\,M$_{\\odot}\\leq M\\leq$\\,0.2\\,M$_{\\odot}$) in Chamaeleon\\,I exhibit 6.7\\,$\\mu$m excess emission but no K-band excess. It thus seems possible that towards less massive objects one would expect to find comparatively fewer objects with near-infrared excess and perhaps none among brown dwarfs. However, Muench et al.\\ (\\cite{muench01}) find a K$_{\\rm s}$-band excess in 65\\% of Trapezium cluster brown dwarfs, slightly more than the 50\\% K-band excess fraction that Lada et al.\\ (\\cite{lada00}) find in more massive stars (0.1$-$1\\,M$_{\\odot}$). Clearly, the K-band excess frequency and presumably the fraction of low-mass stars and brown dwarfs possessing warm circumstellar material is smaller around $\\sigma$\\,Ori than in the 1\\,Myr old Trapezium cluster. This is true for both the low-mass stars and also for the brown dwarfs. Even though our sample of brown dwarfs is small, the lack of even a single K$_{\\rm s}$-band excess detection (in the $JHK_{\\rm s}$ colour-colour diagram) is significant compared with the 65\\% fraction seen by Muench et al. (\\cite{muench01}) in the Trapezium. Our results are consistent with the low disk fractions of 5$-$12\\%, determined from K- and L-band excesses and H$\\alpha$ emission in the similarly aged Upper Scorpius, NGC\\,2362 and $\\lambda$\\,Ori associations (Walter et al. \\cite{walter94}; Preibisch \\& Zinnecker \\cite{preibisch99}; Haisch et al.\\ \\cite{haisch01}; Dolan \\& Mathieu \\cite{dolan99}, \\cite{dolan01}). We could interpret our results within the framework of the decreasing disk fraction with age that has been clearly established by Haisch et al.\\ (\\cite{haisch01} and references therein), using the ($K-L$) excesses found among the low-mass stars of several young clusters and associations. Our results would then support their conclusions, but also extend the range of masses over which the disk decay has been observed to the substellar regime. In summary, we have found no evidence for a significant population of low-mass stars and brown dwarfs around $\\sigma$\\,Ori which exhibit K$_{\\rm s}$-band excesses. This points strongly towards a low disk frequency, consistent with observations of other similarly aged clusters. Thus this new result adds support to the idea that disk evolution in clusters takes place on short (few Myr) timescales and extends this conclusion to very low masses. We find some evidence to suggest that disk frequencies in the $\\sigma$\\,Ori association, as measured by the K-band excess, may be mass dependent in the opposite sense to those found in the Trapezium cluster. We do not attach great weight to this result at present because of difficulties in comparing disk fractions measured by different authors using differing techniques but, if true, it might suggest that the disk dissipation timescales were shorter for lower mass objects. More answers may be obtained on the mass dependence of the disk frequency by extending infrared surveys into the brown dwarf domain for more clusters with different ages, and preferably at longer wavelengths, where disk presence can be more reliably diagnosed." + }, + "0112/astro-ph0112405_arXiv.txt": { + "abstract": "The nucleosynthesis and production of radioactive elements in SN 1987A are reviewed. Different methods for estimating the masses of \\nia, \\nib, and \\ti~ are discussed, and we conclude that broad band photometry in combination with time-dependent models for the light curve gives the most reliable estimates. ", + "introduction": "\\label{sec_introd} Two of the outstanding issues for SNe are the nucleosynthesis and the explosion mechanism. Although SNe have been known from the 1950's to be the most important sources of heavy elements, there is little quantitative evidence for this. SN 1987A has in this respect been a unique source of information thanks to the possibility to obtain spectral information at very late epochs when the ejecta is transparent and the central regions observable. The nature of the explosion mechanism is still to a large extent unknown, and constraints from observations are badly needed. The main such constraints are provided by the hydrodynamic structure, e.g., the extent of mixing, and the masses of the iron peak elements, in particular the radioactive isotopes \\citep[e.g.,][]{Timmes et al. (1996), Kumagai et al. (1993), 2000ApJ...531L.123K}. ", + "conclusions": "SN 1987A is by far the best source of nucleosynthesis of any SN. Reliable masses for both the most abundant elements, H, He, and O, as well as the the most abundant radioactive isotopes, \\nia, \\nib, and \\ti~ have been obtained. These provide some of the most important constraints on the progenitor and on the explosion itself. CF is grateful to Roland Diehl for organizing a very successful workshop, and to Don Clayton and Nikos Prantzos for useful discussions and comments." + }, + "0112/astro-ph0112363_arXiv.txt": { + "abstract": "The discovery of hard X-ray sources near the top of a flaring loop by the HXT instrument on board the YOHKOH satellite represents a significant progress towards the understanding of the basic processes driving solar flares. In this paper we extend the previous study of limb flares by Masuda (1994) by including all YOHKOH observations up through August 1998. We report that from October 1991 to August 1998, YOHKOH observed 20 X-ray bright limb flares (where we use the same selection criteria as Masuda), of which we have sufficient data to analyze 18 events, including 8 previously unanalyzed flares. Of these 18 events, 15 show detectable impulsive looptop emission. Considering that the finite dynamic range (about a decade) of the detection introduces a strong bias against observing comparatively weak looptop sources, we conclude that looptop emission is a common feature of all flares. We summarize the observations of the footpoint to looptop flux ratio and the spectral indices. We present light curves and images of all the important newly analyzed limb flares. Whenever possible we present results for individual pulses in multipeak flares and for different loops for multiloop flares. We then discuss the statistics of the fluxes and spectral indices of the looptop and footpoint sources taking into account observational selection biases. The importance of these observations (and those expected from the scheduled HESSI satellite with its superior angular spectral and temporal resolution) in constraining acceleration models and parameters is discussed briefly. {\\it Subject Headings:} Sun:flares--Sun:X-rays--acceleration of particles ", + "introduction": "The discovery of hard X-ray sources located at or above the top of solar flare loops by the HXT instrument on board the YOHKOH satellite has provided great insight into the processes that drive solar flares. The canonical event for this phenomenon is the flare of January 13, 1992, first described by Masuda et al. (1994), and later analyzed by Alexander \\& Metcalf (1997), and is commonly referred to as the ``Masuda'' flare. This flare, which is clearly delineated by a soft X-ray (thermal) loop, shows three compact hard X-ray sources, two located at the footpoints (FPs for short) and a third above the top of the soft X-ray loop. The first systematic study of such sources was undertaken by Masuda (1994). During the period of time between the satellite's first scientific observations (October 1991) up to September 1993, Masuda selected 11 X-ray bright limb flares observed by YOHKOH that met his selection criteria (see below), one of these 11 events was interrupted event and was not analyzed further. Of the remaining 10 flares, 6 demonstrated a nonthermal looptop (LT for short) source while another demonstrated what Masuda termed a ``super hot'' thermal LT source. This indicates that LT hard X-ray emission is fairly common. This view is strengthened further considering the limited dynamic range of the HXT instrument and the result we describe in this paper. Thus, it is reasonable to conclude that LT sources are present in all flares. It is generally agreed that the primary energy of solar flares must come from reconnection of stressed magnetic fields, and as pointed out by Masuda et al. (1994), the YOHKOH LT observations lend support to theories that place the location of energy release high up in the corona. The energy released by reconnection can be used to heat the ambient plasma and/or to accelerate electrons and protons to relativistic energies. The power law hard X-ray spectra seen in many of the LT sources indicate that nonthermal processes, such as particle acceleration, are indeed occurring at or near these locations. The exact mechanism of the acceleration is a matter of considerable debate. In previous works (see Petrosian 1994 and 1996) we have argued that among the three proposed particle acceleration mechanisms (electric fields, shocks, and plasma turbulence or waves) the stochastic acceleration of ambient plasma particles by plasma waves provides the most natural means of explaining the observed spectral and spatial features of solar flares. In two recent works (Petrosian \\& Donaghy, 1999 and 2000, hereafter {\\bf PD}) we demonstrated how the power-law spectral indices and emission ratios (obtained by Masuda 1994) can be used to constrain the model parameters. In this paper, we expand and extend Masuda's analysis for further investigation of the ubiquity and nature of the LT source, and to gain a clearer picture of the range of values of the fluxes and spectral indices of the FP and LT sources. Furthermore, we investigate the temporal evolutions of the images of several new flares observed by HXT to determine the relationship between the many spatially and temporally distinct sources that occur in complex flaring events. In the next section we describe our procedure, and in \\S 3 we present the light curves and images of most of the new events, including three events which appear to be examples of multiple loop flares. In \\S 3 present the statistics of the relative fluxes and spectral indices of the LT and FP sources. In \\S 5 we present a summary and our conclusions. ", + "conclusions": " 1. The LT hard X-ray emission seems to be a common characteristic of the impulsive phase of solar flares, appearing in some form in 15 of the 18 selected flares. The absence of LT emission in the remaining cases is most likely due to the finite dynamic range of the imaging technique. 2. The ratio of the summed FP to LT fluxes, lies in the range $10 \\gtrsim {\\cal R}=FPs/LT \\gtrsim 1$ and has a relatively flat distribution in this range. The lower limit is intrinsic to the process but the upper limit is most likely due to the finite (one decade) dynamic range of the imaging technique. Because of this it is difficult to know the true distribution of this ratio. 3. The overall distribution of the power-law spectral indices rises rapidly above $\\gamma=2$, peaks around 4 and declines gradually thereafter. This is similar to previous determinations of this distributions from HXRBS on board the {\\it Solar Maximum Mission} (McTiernan \\& Petrosian 1991), but contains a few more steep spectra, specially for LT sources, which could be due to thermal contamination and lower energy range of the HXT compared to HXRBS. This agreement is encouraging and indicates the reliability of our spectral determination. 4. The spectra tend to steepen at higher energies (spectral index $\\gamma$ increases by 1 to 2), especially for sources with $\\gamma <6$, for which the thermal contribution should be the lowest. This is the opposite of what is observed at higher energies, where spectra tend to flatten above 100's of keV (McTiernan \\& Petrosian 1991). The directivity of the X-ray emission and the albedo effect for the limb flares under consideration could also play some role, especially for the FP source (in the model discussed by {\\bf PD} the LT source should emit isotropically). However, this is expected to be a small effect at low energies ($<100$ keV) under consideration here. 5. The spectral index $\\gamma$ of LT sources is larger ({\\it i.e.} spectra are steeper) than that of the FP sources on the average by one (${\\bar\\gamma}_{LT}=6.2\\pm 1.5, {\\bar\\gamma}_{FP}=4.9\\pm 1.5$). The differences between directivity and the albedo effects mentioned above could be a partial reason for such a behavior, but the physics of the acceleration process must certainly play a role here. 6. We have also described how the above results can be used to constrain the model parameters especially those related to the acceleration process. For example, for models with acceleration at the LT source (see, e.g. Park, Petrosian \\& Schwartz 1997 and {\\bf PD}), the observed ranges of the flux ratios and spectral indices indicate a rapid escape of the accelerated electron relative to the acceleration timescale, which is related to the momentum diffusion coefficient in the acceleration process. Moreover, the difference between spectral indices of LT and FP sources can constrain the energy dependence of the escape time which is related to the pitch angle diffusion rate in the acceleration site. This demonstrates that similar studies of these characteristics of the flares can yield important information about the genesis and evolution of solar flares, and we eagerly anticipate the increased spectral, temporal and spatial resolution possible with the instruments of the upcoming HESSI satellite. Finally, we note that solar flares occur in many different morphologies, the most common being a simple flaring loop with one LT and two FP sources. However, as we discussed in \\S 3.2, interacting loop models and even more complicated structures are frequently observed. We would like to thank an anonymous referee for a careful reading of the original manuscript and for numerous helpful comments and suggestions that improved the paper considerably. This work is supported in parts by NASA grants NAG-5-7144-0002 and NAG5-8600-0001 and by a fellowship to TQD from Stanford's Undergraduate Research Opportunities. \\newpage" + }, + "0112/astro-ph0112155_arXiv.txt": { + "abstract": "We report on Australia Telescope Compact Array (ATCA) \\HI\\ observations carried out in the direction of bilateral supernova remnants (SNRs) with associated neutron stars: G296.5+10.0 and G320.4--1.2, in a search for the origin of such morphology. From these studies we conclude that in the case of G296.5+10.0, located far from the Galactic plane, the \\HI\\ distribution has not influenced the present morphology of the SNR. In the case of G320.4--1.2, evolving in a denser medium, the combined action of the central pulsar, PSR B1509-58, with the peculiar distribution of the surrounding medium, has determined the observed characteristics of the SNR. ", + "introduction": "Several Galactic supernova remnants (SNRs) exhibit an unusual bilateral morphology, characterized by a clear axis of symmetry, two bright limbs on either side and low level of emission near the top and bottom along the symmetry axis. The origin of this ``barrel-shaped'' appearance has provoked considerable debate for the past few years. A detailed study of the gaseous environs of bilateral SNRs is a very useful tool to disentangle intrinsic origins (like the presence of biconical beams from a central neutron star (NS), asymmetric explosions, etc.) from extrinsic causes (stratification of the interstellar density, the strength and orientation of the ambient magnetic field, etc.). As a part of an ongoing project to observe the environs of bilateral SNRs, we have conducted a detailed \\HI\\ study around G296.5+10.0 and G320.4--1.2. These two bilobular SNRs share the characteristic of harbouring an eccentric X-ray pulsar in the interior (the source 1E 1207.4--5209 inside G296.5+10.0, Helfand \\& Becker 1984, and PSR B1509-58 associated with G320.4--1.2, Seward \\& Harnden 1982). Using the Australia Telescope Compact Array (ATCA) we surveyed wide fields around these extended SNRs, looking for peculiar alignments, and/or properties of the surrounding gas which may give clues as to the origin of the observed radio morphology. ", + "conclusions": "" + }, + "0112/astro-ph0112225_arXiv.txt": { + "abstract": "{We present optical $V$ and $i$-band light curves of the gravitationally lensed BAL quasar HE~2149$-$2745. The data, obtained with the 1.5m Danish Telescope (ESO-La Silla) between October 1998 and December 2000, are the first from a long-term project aimed at monitoring selected lensed quasars in the Southern Hemisphere. A time delay of $103\\pm12$ days is determined from the light curves. In addition, VLT/FORS1 spectra of HE~2149$-$2745 are deconvolved in order to obtain the spectrum of the faint lensing galaxy, free of any contamination by the bright nearby two quasar images. By cross-correlating the spectrum with galaxy-templates we obtain a tentative redshift estimate of $z=0.495\\pm0.01$. Adopting this redshift, a $\\Omega=0.3$, $\\Lambda=0.7$ cosmology, and a chosen analytical lens model, our time-delay measurement yields a Hubble constant of H$_{0}=66 \\pm 8\\, {\\rm km}~{\\rm s^{-1}}~{\\rm Mpc^{-1}}$ (1$\\sigma$ error) with an estimated systematic error of $\\pm3\\, {\\rm km}~{\\rm s^{-1}}~{\\rm Mpc^{-1}}$. Using non-parametric models yields H$_{0}=65 \\pm 8\\, {\\rm km}~{\\rm s^{-1}}~{\\rm Mpc^{-1}}$ (1$\\sigma$ error) and confirms that the lens exhibits a very dense/concentrated mass profile. Finally, we note, as in other cases, that the flux ratio between the two quasar components is wavelength dependent. While the flux ratio in the broad emission lines - equal to 3.7 - remains constant with wavelength, the continuum of the brighter component is bluer. Although the data do not rule out extinction of one quasar image relative to the other as a possible explanation, the effect could also be produced by differential microlensing by stars in the lensing galaxy. ", + "introduction": "The time-delay between the gravitationally lensed images of a distant source is a measurable parameter. Observed as the time difference between the arrival dates of a single (lensed) wavefront emitted by a distant source, it is directly related to the Hubble constant H$_{0}$ (Refsdal~\\cite{Refsdal}). Obtaining accurate time-delay measurements in multiply lensed quasars can therefore yield {\\bf (i)} a determination of H$_{0}$ provided the mass distribution in the lens is known, or {\\bf (ii)} constraints on the mass distribution in a given lens, using H$_{0}$ as inferred from other methods. During the last 20 years much effort has been devoted to the observations of lensed quasars, and in particular to long-term monitoring of selected systems. Some of these are Q0957+561 (Schild \\cite{Schild}, Vanderriest et al. \\cite{Vanderriest}), PG1115+080 (Schechter et al. \\cite{Schechter}), B1608+656 (Fassnacht et al. \\cite{Fassnacht}) and B1600+434 (Burud et al. \\cite{Burud} and Koopmans et al. \\cite{Koopmans}). In this context, we have been conducting a photometric monitoring program at the Danish 1.5-m telescope at La Silla observatory (ESO, Chile) since October 1998, with the goal of measuring the time-delays in several well studied lensed quasars. We present here the first result from this program: the time-delay measurement in the two-image quasar \\object{HE~2149$-$2745}. The lensed nature of the BAL quasar \\object{HE~2149$-$2745} at $z=2.03$ was established by Wisotzki et al. (\\cite{Wisotzki}). This system proves to be an easy target for monitoring at a site with reasonable seeing conditions (up to 2\\arcsec): it is bright ($B=17.3$) and the two quasar images have an angular separation of 1.7\\arcsec. The monitoring program, the light curves and the time delay are discussed in Sect.~\\ref{sect:data}, ~\\ref{sect:phot} and ~\\ref{sect:timedelay} below. The lensing galaxy has been detected in HST NICMOS and WFPC2 images but its redshift remains unknown. With the aim of measuring this redshift, we have obtained a spectrum of \\object{HE~2149$-$2745} with FORS1 at UT1 (ESO-Paranal, Chile). The analysis of the spectroscopic data is described in Sect.~\\ref{sect:spectro}. This section also includes a discussion on the spectral differences between the two quasar components. Mass models and estimates of the Hubble constant are presented in Sect.~\\ref{sect:mass}. Finally, Sect.~\\ref{sect:discussion} summarises the main results. \\begin{figure} \\centering \\caption{Field of view 5\\arcmin x5\\arcmin\\, in size around \\object{HE~2149$-$2745}. The three reference stars (labelled S1, S2, S3) used for the photometry and the two PSF-stars (PSF1, PSF2) used in the spectral deconvolution are indicated. North is up and East to the left.} \\label{field} \\end{figure} ", + "conclusions": "\\label{sect:discussion} We have presented the first result of a long-term photometric monitoring campaign undertaken at ESO between 1998 and 2000 at the 1.54-m Danish telescope. Our $V$ and $i$ light curves allow us to measure a time-delay of $\\Delta t = 103 \\pm12$ days between the two quasar images of \\object{HE~2149$-$2745}. From VLT spectroscopy, we have derived a tentative estimate of the lens redshift to be $z=0.489$. Applying both analytic and numerical lens models to the case of \\object{HE~2149$-$2745}, we derive H$_0 = 66 \\pm 6\\, \\ho$ with an additional systematic error of $\\pm\\, 3\\, \\ho$ in the case of the analytic models, due the limited range of the ellipticity. The derived mass models are relatively compact, although not as compact as the light profile of the galaxy. An extra source of systematics might be introduced by the uncertainty on the lens redshift estimate. This is however not critical for the determination of H$_0$ as the error is dominated by the uncertainty in the time-delay measurement. As \\object{HE~2149$-$2745} shows smooth light curves, it is likely that the situation can be improved by continued monitoring. With an improved time-delay it would also be highly desirable to re-determine the lens redshift more precisely. Our monitoring program of \\object{HE~2149$-$2745} is the first to be carried out in two bands on such a regular basis and over such a long time scale. Given the error bars, we do not see any significant colour variation over the 900 days of observation. Moreover, our spectra of the two quasar images show that the flux ratios in the broad emission lines behave differently for the continuum flux ratio, and that the flux ratio measured in the BAL structure of the source follows the behavior as the continuum region. Such behavior can be explained both by microlensing or by differential extinction by the lensing galaxy, or both. So far, the data do not allow us to distinguish between the two possible explanations. In order to confirm microlensing, one would need for example to know the time-scale of the putative event, i.e., to measure the absolute magnification of the event and its duration. This would not only allow us to confirm microlensing, but also to use it to map the radial structure of the central AGN in the source. Although the Einstein radius crossing time is long for \\object{HE~2149$-$2745}, of the order of 10 years, much shorter microlensing events, such as a caustic crossings, can occur. Conducting a long-term spectrophotometric monitoring could therefore allow us to probe the AGN size \\object{HE~2149$-$2745}." + }, + "0112/astro-ph0112539_arXiv.txt": { + "abstract": "We present models for the complete life and death of a 60\\,\\Msun star evolving in a close binary system, from the main sequence phase to the formation of a compact remnant and fallback of supernova debris. After core hydrogen exhaustion, the star expands, loses most of its envelope by Roche lobe overflow, and becomes a Wolf-Rayet star. We study its post-mass transfer evolution as a function of the Wolf-Rayet wind mass loss rate (which is currently not well constrained and will probably vary with initial metallicity of the star). Varying this mass loss rate by a factor 6 leads to stellar masses at collapse that range from 3.1\\,\\Msun up to 10.7\\,\\Msun. Due to different carbon abundances left by core helium burning, and non-monotonic effects of the late shell burning stages as function of the stellar mass, we find that, although the iron core masses at collapse are generally larger for stars with larger final masses, they do not depend monotonically on the final stellar mass or even the C/O-core mass. We then compute the evolution of all models through collapse and bounce. The results range from strong supernova explosions ($\\Ekin > \\Ep{51}\\,\\erg$) for the lower final masses to the direct collapse of the star into a black hole for the largest final mass. Correspondingly, the final remnant masses, which were computed by following the supernova evolution and fallback of material for a time scale of about one year, are between 1.2\\,\\Msun and 10\\,\\Msun. We discuss the remaining uncertainties of this result and outline the consequences of our results for the understanding of the progenitor evolution of X-ray binaries and gamma-ray burst models. ", + "introduction": "\\lSect{intro} It has long been known that stellar-mass black holes could form from the collapse of massive stars (Oppenheimer \\& Snyder 1939) and it is believed that above some progenitor initial mass limit, stars collapse to form black holes. A growing set of evidence suggests that this mass limit in single stars ($M_{\\rm BH}^{\\rm S}$) lies somewhere below 25\\,M\\sun. Fryer (1999) obtained this result from core collapse simulations. Maeder (1992) and Kobulnicky \\& Skillman (1997) find it to be most consistent with nucleosynthesis constraints. And the explosion of a star of about 20\\,M\\sun as SN~1987A (Arnett et al. 1989) gives a lower limit to the initial stellar mass required for direct black hole formation (to be distinguished from black hole formation due to fall back; see below). Within the uncertainties of the above studies, it appears that the black hole mass limit of single stars is reasonably constrained. In this paper, we deal with the black hole mass limit for primary stars of close binary systems ($M_{\\rm BH}^{\\rm B}$), i.e., for stars which evolve into the compact objects contained in high and low mass X-ray binaries or which might become $\\gamma$-ray burst sources: collapsars and black hole binary mergers (see Fryer, Woosley, \\& Hartmann 1999 for a review). As the presence of a companion star only increases the mass loss of the primary star in a binary system before the first supernova, the black hole mass limit of binaries is likely to be larger than that of single stars (i.e., $M_{\\rm BH}^{\\rm B} > M_{\\rm BH}^{\\rm S}$). There are two observational constraints from observed black hole systems on the black hole mass limit in binaries. Using population synthesis studies, Portegies Zwart, Verbunt, \\& Ergma (1997) find that the number of low mass black hole X-ray binaries in our Galaxy requires $M_{\\rm BH}^{\\rm B} < 25\\,\\Msun$. Ergma \\& van den Heuvel (1998) argue (without detailed modelling) that the observed periods of less than $\\sim 10\\,$h found in most low mass black hole systems are incompatible with $M_{\\rm BH}^{\\rm B} > 25\\,\\Msun$. While the latter argument is based on considerations of angular momentum loss associated with the (uncertain!) Wolf-Rayet winds. In both investigations neither $M_{\\rm BH}^{\\rm S}$ and $M_{\\rm BH}^{\\rm B}$ are distinguished, nor do they consider the dependence of $M_{\\rm BH}^{\\rm B}$ on the type of binary evolution. Wellstein \\& Langer (1999; WL99) showed that indeed $M_{\\rm BH}^{\\rm B}$ may be very different from $M_{\\rm BH}^{\\rm S}$, and is strongly dependent on the type of binary interaction. Only in the so called Case~C systems --- i.e., in initially wide systems where mass transfer starts only after the primary has evolved through the major part of core helium burning --- may the black hole mass limits be comparable; i.e., $M_{\\rm BH}^{\\rm BC} \\simeq M_{\\rm BH}^{\\rm S}$ (Brown, Lee, \\& Bethe 1999). However, due to the Wolf-Rayet winds, which reduce the total stellar mass and thus the helium core mass {\\em during} core helium burning, the black hole mass limit for Case~A and Case~B binaries (mass transfer starts during or directly after core hydrogen burning), is clearly smaller; i.e., $M_{\\rm BH}^{\\rm BA} > M_{\\rm BH}^{\\rm BC}$ and $M_{\\rm BH}^{\\rm BB} > M_{\\rm BH}^{\\rm BC}$. And although WL99 found that $M_{\\rm BH}^{\\rm BA} \\simeq M_{\\rm BH}^{\\rm BB}$, this result is expected to be different (i.e., $M_{\\rm BH}^{\\rm BA} > M_{\\rm BH}^{\\rm BB}$) if less efficient Wolf-Rayet wind mass loss were assumed (i.e., compare our Model 1s2 below with Model 2$^{\\prime}$ of WL99). The relevance of detailed binary evolution models for the significance of constraints on $M_{\\rm BH}^{\\rm B}$ from observed X-ray binaries has been demonstrated by WL99. Ergma \\& van den Heuvel (1998) argued that the pulsar GX301-2, which has a 40...50\\,M\\sun supergiant companion, originated from a star of more than 50\\,M\\sun, implying that $M_{\\rm BH}^{\\rm B} > 50 M_{\\odot}$ for this particular system. However, WL99 computed detailed Case~A progenitor evolution models which satisfy all observational constraints for this system, but in which the pulsar progenitor has an initial mass of only 26\\,M\\sun. This solution implies $M_{\\rm BH}^{\\rm S} < M_{\\rm BH}^{\\rm BA} < 26 M_{\\odot}$. In the following we investigate in detail the black hole mass limit for primary stars in close binary systems, using a particular Case~B evolutionary sequence with 60\\,M\\sun primary star. We study five cases where we only change the stellar wind mass loss rate in the Wolf-Rayet phase of the primary star which we vary over the anticipated regime of uncertainty, i.e., by a factor of 6. For high mass loss rates (e.g., Hamann, Koesterke \\& Wessolowski 1995, Langer 1989b, Braun 1997, Woosley, Langer, \\& Weaver 1995), the stellar mass at collapse can reduce down to $\\sim$3\\,M\\sun with an iron core mass as low as 1.3\\,M\\sun (Woosley, Langer, \\& Weaver 1995). The lower Wolf-Rayet mass loss rates proposed by Hamann \\& Koesterke (1998) and Nugis \\& Lamers (2000) result in higher final masses (WL99). For all five cases, we model the evolution of the primary star up to iron core collapse (Sect.~2), which predicts the final stellar mass and the detailed pre-supernova structure of these objects --- e.g., their final iron core mass. We then use these pre-collapse structures to model their collapse through bounce and explosion (if applicable; Sect.~3). Finally, we follow the evolution of the supernova ejecta for about one year, which allows us to estimate the final remnant mass including fall back of previously ejected material onto the compact object. With these coupled simulations we obtain the dependence of the compact remnant mass in a given binary system on the assumed strength of the Wolf-Rayet wind. Implications of our results for determinations of the black hole mass limit and the progenitor evolution of X-ray binaries are discussed in Sect.~4. ", + "conclusions": "In this paper, we present models for life and death of a series of 60\\,\\Msun stars in binaries, whose evolution differs only by the mass loss rate adopted in the Wolf-Rayet stage. The Wolf-Rayet mass loss rate is still very uncertain and has only recently been revised downward by a factor 2$\\ldots$3 (Hamann \\& Koesterke 1998). In addition to being uncertain, it is likely that the Wolf-Rayet mass-loss rate depends on metallicity. Aside from the mass-loss rate, the metallicity has very little effect on the stellar evolution models. Thus, the study presented here can also be regarded as a study of this binary at different metallicities -- the uncertainty of the WR mass loss rate then translates into a variation of the the initial stellar metal abundance. We obtain final pre-collapse stellar masses in the range from 3.1\\,\\Msun for the highest mass loss rate (Model~1s1) up to 10.7\\,\\Msun for the lowest mass loss rate (1s6), while the central carbon abundance at core helium exhaustion drops form 35\\,\\% to 22\\,\\%, respectively. The ensuing complex interaction of carbon core and shell burning phases with later burning stages causes a non-monotonic behavior of the pre-collapse structure of the star, i.e., the masses of the neon-oxygen core, the silicon core or the deleptonized core (\\Tab{evo}; \\Sect{Cc2cc}). Following the subsequent collapse and supernova explosion of these stars, we find that these models produce a range of compact remnants from a 1.17\\,\\Msun neutron star (1s2) to a 10.7\\,\\Msun black hole (1s6). The remnant mass does not scale strictly with the iron core mass (it depends on the density and temperature structure of the collapsing core - see \\S 3) and we can not use this core mass to estimate the supernova explosion energy or the compact remnant mass. Amazingly, the large differences in the remnant mass are caused by only a factor of~ 3 change in the Wolf-Rayet mass-loss rate, and a 40\\% difference in the amount of mass lost through Wolf-Rayet winds. In view of this extreme sensitivity of the remnant mass on the WR mass loss, the persisting uncertainty of the WR mass loss rate, and the uncertainties of our core collapse models, it is difficult to draw solid conclusion. Nevertheless, taking our results at face value implies that to form a black hole of 10\\,\\Msun or more from a 60\\,\\Msun star in a Case~B (or Case~A) binary, one might need to use WR mass loss rates smaller than the currently favoured ones (unless the system is at lower metallicity). That is, it seems difficult (though not impossible) to form those X-ray binaries which contain the most massive black holes, like Cyg~X1 or V404~Cyg, through Case~A or~B at solar metallicity, and the Case~C scenario (Brown et al. 1999, 2001, Kalogera 2001, Fryer \\& Kalogera 2001) may provide a viable alternative. For most low mass black hole binaries, which may contain black holes with 3$\\ldots$7\\,\\Msun (Fryer \\& Kalogera 2001), our results imply that a Case~B progenitor evolution may be sufficient if we assumed that the mass loss rate were reduced by a factor of $\\sim$4. Given the uncertainties in WR mass loss, stellar evolution, and core-collapse, Case~B progenitors for most black hole binaries are not excluded. Furthermore, only with our lowest mass loss rate were we able to produce a direct collapse black hole from this Case~B progenitor. Fryer, Woosley, \\& Hartmann (1999) suggested that most collapsar $\\gamma$-ray burst progenitors are produced in binares that undergo Case~B mass transfer, but our results imply that such binaries may not produce collapsars (at least at solar metallicity). However, most $\\gamma$-ray bursts occur at high redshifts and low metallicities. That is, if the Wolf-Rayet mass loss rates truly decrease with decreasing metallicity, this does not preclude Case~B progenitors of $\\gamma$-ray bursts. All in all, our results should be understood as temporary, awaiting a better understanding of the Wolf-Rayet winds and the core collapse of massive stars. But to understand the origin of the black hole binaries in our nearby universe, we must continue to pursue Case AB, as well as Case~C models. We should finally mention that we left out two potentially important stellar parameter which is essential for most current models of collapsing stars: rotation and magnetic fields. Rotation may add another dimension to the expected remnant mass as function of stellar parameters, which has to be left here for future investigations. Although Fryer \\& Heger (2000) found that the currently proposed dynamos would not develop strong enough magnetic fields to drive the explosion alone, magnetic fields may still effect the explosion and should be considered." + }, + "0112/hep-ph0112018_arXiv.txt": { + "abstract": "DAMA is searching for rare processes by developing and using several kinds of radiopure scintillators: in particular, NaI(Tl), liquid Xenon and CaF$_2$(Eu). The main results are here summarized with particular attention to the investigation of the WIMP annual modulation signature. ", + "introduction": "DAMA is devoted to the search for rare processes by developing and using low radioactivity scintillators. Its main aim is the search for relic particles (WIMPs: Weakly Interacting Massive Particles). In addition, due to the radiopurity of the used detectors and of the installations, several searches for other possible rare processes are also carried out, such as e.g. on exotics, on $\\beta\\beta$ processes, on charge-non-conserving processes, on Pauli exclusion principle violating processes, on nucleon instability and on solar axions\\cite{Bel96,ncim,Bepa,Ca1,Ca2,astr,Becha,Beel,Rd99,Pl99,Ndn00,Beax,Xe136}. The main experimental set-ups running at present are: the $\\simeq$ 100 kg NaI(Tl) set-up, the $\\simeq$ 6.5 kg liquid Xenon (LXe) set-up and the so-called ``R\\&D'' apparatus. Moreover, a low-background germanium detector is operative underground for measurements and selections of sample materials. In the following the most recent results achieved with the $\\simeq$ 6.5 kg LXe set-up and with the ``R\\&D'' apparatus will be briefly recalled, while a more dedicated discussion will be devoted to the present results on the investigation of the WIMP annual modulation signature by the $\\simeq$ 100 kg NaI(Tl) set-up. ", + "conclusions": "" + }, + "0112/astro-ph0112007_arXiv.txt": { + "abstract": "We show that neutrinos can be produced through standard electroweak interactions in matter with time-dependent density. ", + "introduction": "In this talk I will show that neutrinos can be pair produced by a time-dependent matter density through standard electroweak interactions. The idea is that both ordinary matter and the nuclear matter found in neutron stars carry a net SU(2) charge. Neutrinos couple to this charge through the electroweak interactions. In a time-dependent matter background this leads to neutrino pair production~\\cite{nu_prod}. This phenomenon is analagous to electron pair production by a time-varying electromagnetic field~\\cite{ct,grib}, or fermion production by a variable scalar field as can occur during preheating~\\cite{baa,gk}. Applications include density waves in a neutron star, a neutron star in a binary system, or a supernova explosion. We expect the amount of neutrinos produced to be small though. ", + "conclusions": "To summarize, neutrinos can be pair produced by a time-dependent matter background. Although neat, the phenomenon seems of little practical importance. Due to the smallness of all mass scales involved, the amount of neutrinos produced is small." + }, + "0112/astro-ph0112188_arXiv.txt": { + "abstract": " ", + "introduction": "A variety of observations make a compelling case for the existence of dark matter in unknown form(s). Among the discussed possibilities we may mention axions, WIMPS, cosmic strings, brown dwarfs and Primordial Black Holes, see \\cite{[PEA98]}. It has been longely recognized that the abundance of evaporating PBHs \\cite{[HAW75]} may be constrained by a variety of methods . For example, Carr and Mac Gibbon (see \\cite{[CMC98]} and references therein) used the gamma-ray diffuse background to establish constraints. Similarly, Liddle and Green \\cite{[LG98]} reviewed a wide variety of methods to study a possible PBH population. Generally speaking, the methods dealing with the evaporation of PBHs apply to the final phase of these objects only, and therefore do not make full use of the radiative thermal flux at larger wavelengths predicted by the Hawking process of quantum evaporation \\cite{[CMC98]},\\cite{[LG98]}. The formation of PBHs from primordial fluctuations was first studied by Carr \\cite{[CARR75]}. In that work he established that large, gaussian primordial fluctuations from a scale-free primordial spectrum may be responsible for a power-law PBH mass function. A new twist to the problem has been added by the recent work of Niemeyer and Jedamzik \\cite{[NJ98]},\\cite{[NJ99]} who used the concepts of critical phenomena to study the initial mass function for Gaussian and \"mexican hat\" initial fluctuation spectra. They have argued that in these cases PBHs may form with masses below the horizon mass, and found a mass function which is not a pure power-law. Black holes may form at all epochs (see Yokoyama \\cite{[YOK98]}, and Kawasaki and Yanagida \\cite{[KY98]} for recent work on PBH formation from the collapse of large density perturbations in double-inflation and supergravity models) and since there is considerable uncertainty about the actual processes that effectively formed PBHs, we seek here a general method to constrain the population which may be used independently of the specific PBH formation mechanism(s) (see for instance, Hansen et. al., and Bousso \\cite{[BOU98]} for a discussion of other possibilities). We shall evaluate a simple (but useful) model consisting in a single-mass scale, which may be associated to the typical mass-scale of any given mass function showing such a peak (this situation is very well satisfied by models in which large perturbations collapse to form PBHs). The simplest reexamination of PBH physics has shown that, as a result of particle absorption from the expanding background, there is a maximum value of the mass (termed \"critical mass\" in \\cite{[CH98]}) splitting an evaporating (or subcritical) from a non-evaporating (or supercritical) subset of a general mass distribution. The critical mass separating these regimes for each redshift is therefore a useful boundary to discuss the physics of the PBHs and will be defined in the next section . A detailed quantitative discussion of these issues has been given in Refs.\\cite{[CH98]},\\cite{[CH99]}. Due to the absorption of particles from the thermal background, PBH evaporation (and therefore some limits to their abundance) have to be reexamined. Actually we shall see that the existence of a regime of $\\sim$ constant mass in the life of PBHs is very important for a full evaluation of their survival. Previous studies \\cite{[BAR91]} have not considered the regime of quasi-constant mass and thus need to be updated to account for the evaporation at each evaporation scale labeled by the redshift. ", + "conclusions": "\\bigskip We have shown that under a set of circumstances a limit on $\\Omega_{pbh}$ better than ${10^{-10}}$ can be obtained by using the background brightness of the sky. This is quite stringent if we compare it with the values given by other methods \\cite{[LG98]}. We attempted to evaluate the evolution of maximal abundance $\\Omega_{pbh}(z)$ for any $z$ from observations performed today ($z=0$), a process that requires integration over $z$ (and eventually over the actual mass spectrum from whatever the formation process, not attempted here). The good news here is that PBHs can be excluded in a fairly large window (as seen in Fig. 4) if they were born above the Hawking mass (but below $10^{26} \\, g$). We see four physical reasons for this fact: 1) Many of these PBHs {\\it still} evaporate today, and therefore contribute to the background in various bands. 2) The cosmic radiation today is much less intense than at high redshift values, and therefore constitutes a useful natural background. 3) The cosmic expansion today is quite weak, then the Doppler damping of the emitted radiation does not have any killing effect. 4) PBHs initially above the critical mass at some $z_{f}$ delay the start of their evaporation, therefore they injected energy later, when the expansion is less effective to damp the radiation. The constraints so obtained are weakly dependent of $q_{0}$ or the cosmological constant, and the mass range probed by this method is larger than previous works. Even though we take into account the existence of the critical mass and the redshift of the emitted radiation through the function $I(\\epsilon_{\\odot},z_{f})$, we found that our analysis renders useful limits for PBH masses initially above $\\sim 10^{15} \\, g$, but not for those below it. In addition, PBHs above ${10}^{26}g$ must be limited by other methods because they do not evaporate at all. The consideration of the different cases leaded us to conclude that $\\epsilon_{\\odot}$ need not to be imposed by hand. Moreover, it happens to be an extremely small number in all cases, with microscopic associated distances $D_{L}(\\epsilon_{\\odot})$. Physically, this means that the spatial distribution of PBHs does not need to satisfy any specific requirement near the earth, and thus the bounds so obtained are quite general. The upper limits to the number density of PBHs were obtained from the simple requirement $\\delta{F}_{pbh} \\, < \\, \\delta{F}_{back}$, since other astrophysical mechanisms may be also contributing to the cosmic background brightness as measured by our detectors. The limits have been derived for a Dirac's delta mass function, whose main features are much simpler than any other choice. We can think this case as a first approach to the general problem, likely accurate for peaky distributions of a more general type. Detailed calculations of the latter remain an interesting problem for future investigations." + }, + "0112/astro-ph0112141_arXiv.txt": { + "abstract": "We assess the global properties of associated \\civ\\gl\\gl1548,1550 absorption lines measured in the spectra of radio-loud quasars drawn from a near-complete, low-frequency selected sample. The observations span restframe \\civ\\ in two redshift ranges --- $0.70.5$; where $S_{\\nu} \\propto \\nu^{-\\alpha}$) as observed between 408\\,MHz and 4.86\\,GHz (see Kapahi et al. 1998). We also define the orientation-dependent ratio, $R$, to be the ratio of core to lobe luminosity measured at a rest-frame frequency of 10\\,GHz (see Kapahi et al. 1998). For consistency with earlier papers we use cosmological parameters $H_0=50$\\,\\kms Mpc$^{-1}$, $q_0 = 0.5$ and $\\Lambda = 0$ throughout. We note that, over the redshift range under consideration, the resulting linear sizes are consistent to within 10\\% of the values using current best estimates for cosmological parameters (namely $\\Omega_m=0.3$, $\\Omega_{\\Lambda}=0.7$, $H_0=70$\\,\\kms Mpc$^{-1}$). ", + "conclusions": "We present a study of \\civ\\ associated absorption in a highly-complete, homogeneous sample of radio-loud quasars. The main results are: \\begin{itemize} \\item We confirm that absorption is more common in steep-spectrum and lobe-dominated quasars, such that the absorbing material lies away from the jet axis in the orientation-dependent unified models. \\item The strength of \\civ\\ absorption decreases with increasing radio-source size. If we assume that the larger sources are older than the smaller (CSS) ones, then we can attribute the decrease in column density to the growth of the radio-source envelope through the ISM of the host galaxy. \\item From the correlation of \\civ\\ absorption strength with optical spectral slope, we conclude that considerable amounts of dust are associated with the absorbing clouds. Consequently we predict that absorbed quasars will be missed preferentially in optically-selected samples, provided similar schemes apply. \\item We find no evidence for changes in the frequency or strength of the absorbers with redshift from $z\\sim 0.7$ to $z\\sim 3$. This lack of cosmic evolution indicates the absorbers are unaffected by gross galaxy evolution, rather they signal a transient phase which is related specifically to the AGN activity. \\item The combination of these results requires that radio sources are triggered in gas-rich, dusty galaxies, such as those immediately following a starburst, and the dust and gas dissipates over the lifetime of the radio source. Thus, the ionization cones may open up with increasing radio-source age. \\end{itemize} \\bigskip" + }, + "0112/astro-ph0112411_arXiv.txt": { + "abstract": "The DJEHUTY project is an intensive effort at the Lawrence Livermore National Laboratory (LLNL) to produce a general purpose 3-D stellar structure and evolution code to study dynamic processes in whole stars. ", + "introduction": "Stellar models in 1-D work remarkably well for most stars. However, stars are three dimensional objects and the computing power is now at a point where we can do better than 1-D models for modeling the large array of physical processes occurring in stars for which spherical symmetry is no longer a valid approximation. With a 3-D stellar code one can tackle the problems linked to rotation, turbulent motions and convection, magnetism, binarity, and explosive phases of stellar evolution in a consistent and physically meaningful way. The DJEHUTY code is an evolution of a radiation hydrodynamics code developed over decades at LLNL. It is our goal to provide the astrophysical community with the first general purpose 3-D stellar structure and evolution code suitable to study the whole gamut of dynamical processes occurring in stars. ", + "conclusions": "" + }, + "0112/astro-ph0112282_arXiv.txt": { + "abstract": "Published parametric models of the Einstein Cross gravitational lens demonstrate that the image geometry can be reproduced by families of models. In particular, the slope of the mass-profile for the lens galaxy is unconstrained. However, recent models predict a dependence of image flux ratios on the slope of the mass profile. We use this dependence to constrain the mass profile by calculating the likelihood of the slope using published mid-IR flux ratios (including microlensing variability). We find that the galaxy is likely to be flatter than isothermal, and therefore that the mass-to-light ratio is decreasing in the inner kpc. ", + "introduction": "Q2237+0305 (The Einstein Cross) was discovered in the CfA Redshift survey (Huchra et al. 1985). The object comprises a quasar at redshift $z=1.695$ that is gravitationally lensed by a foreground galaxy ($z=0.0394$) producing 4 images with separations of $\\sim 1''$. Many models have been proposed for the projected galaxy lens mass distribution based on observations of the lensed images (e.g. Kent \\& Falco 1988, hereafter KF88; Schneider et al. 1988, hereafter S88; Kochanek 1991, hereafter K91; Rix, Schneider \\& Bachall 1992, hereafter RSB92; Wambsganss \\& Paczynski 1994, hereafter WP94; Witt, Mao \\& Schechter 1995, hereafter WMS95; SWL98; Chae, Turnshek \\& Khersonsky 1998, hereafter CTK98). Studies which employ parametric lens models having galaxy and image positions as constraints describe a degeneracy between the ellipticity of the mass profile and its slope $\\nu$. As a result, the total magnification is virtually unconstrained. On the other hand, the models of SWL98 and CTK98 predict variation of the flux ratios with $\\nu$ (this dependence arises due to the slightly asymmetric image geometry). While optical flux ratios cannot be used as model constraints due to microlensing and uncertain differential extinction (e.g. S88; KF88; K91; WP94), measurements in the radio (Falco et al. 1996) should be reliable due to an extended emission region. However Q2237+0305 is faint in the radio, and the observations have large uncertainties. As a result, the majority of published models predict flux ratios consistent with those measured in the radio. Agol, Jones \\& Blaes (2000, hereafter AJB00) observed Q2237+0305 in the mid-IR and measured flux ratios with uncertainties significantly smaller than observations in the radio. Each of the 4 images are observed through the bulge of the lens galaxy which has an optical depth in stars that is of order unity (e.g. KF88; S88; SWL98). This results in a high probability for microlensing, and the mid-IR observations were coincident with significant microlensing of the optical flux. This disparity between optical and mid-IR flux ratios has been used to demonstrate that the mid-IR emission region is large with respect to the microlens Einstein Radius (AJB00; Wyithe, Agol \\& Fluke 2001, hereafter WAF01). The large source size indicates that unlike the optical emission, the mid-IR emission is not subject to significant microlensing. In addition, extinction is not important in the mid-IR. The mid-IR flux ratios should therefore be reliable indicators of the true flux-ratios, and can be used to constrain galaxy mass models. In Sec.~\\ref{macros} we assume microlensing of the mid-IR is negligible and compute the likeli-hood for different published macro-models by comparing the observed mid-IR flux ratios with the predicted flux-ratios. In Sec.~\\ref{MLlim} we discuss the effect of mid-IR microlensing on the results obtained. ", + "conclusions": "Mid-IR flux ratios (Agol et al. 2001) have been used to constrain degeneracies in the slope the mass profile for published models (Chae, Turnshek \\& Khersonsky 1998; Kochanek 1991; Rix, Schneider \\& Bahcall 1992; Schmidt, Webster \\& Lewis 1998; Wambsganss \\& Paczynski 1994) of the bulge of the lensing galaxy G2237+0305. Microlensing of the mid-IR is at a much lower level than for optical emission but may also be important. For the model of Schmidt, Webster \\& Lewis 1998 our calculation includes microlensing of the mid-IR emission by first finding the a-posteriori probability for source size in a macro-model insensitive way using the variation of the image C V-band to mid-IR flux ratio between 2000 (Agol et al. 2001) and 1999 (Agol, Jones \\& Blaes 2000). Microlensing induced flux ratio distributions were produced by combining distributions for the mid-IR source size with magnification distributions corresponding to macro-models having different power law slopes. Likelihoods for the slope of the mass profile were then computed from the combination of mid-IR flux ratio observations (with Gaussian errors) and microlensing induced flux ratio distributions. The mid-IR flux ratios favour a mass profile which is flatter than isothermal, resulting in an ellipticity lower than the light profile. Following the photometric disc-bulge deconvolution of Schmidt (1996), this can be interpreted as a mass-to-light ratio which is higher in the galactic disc than in the bulge. A shallow mass profile results in high magnification, and this helps alleviate the difficulty of reconciling the observed quasar luminosity with microlensing imposed source size constraints (Rauch \\& Blandford 1991). Future mid-IR monitoring with reduced uncertainties will improve constraints on lens galaxy models for 2237+0305 in two ways. Firstly, if no variability of the mid-IR is observed, limits on the mid-IR source size will become tighter, and microlensing a less important consideration. On the other hand, if variability is observed there will no longer be strong dependence of mid-IR source size limits on the upper cutoff of Bayesian prior. Secondly, at the current level of uncertainty, the observations differ by only $\\sim3\\sigma$ from the predicted flux ratios for the least favoured models (from Schmidt, Webster \\& Lewis 1998). Reduced uncertainties will therefore increase the disparity in the likelihoods of different models and result in tighter constraints" + }, + "0112/astro-ph0112447_arXiv.txt": { + "abstract": "Spectra from the Short Wavelength Spectrometer (SWS) on \\iso\\ exhibit artifacts at 4.5 and 8~$\\mu$m. These artifacts appear in spectra from a recent data release, OLP 10.0, as spurious broad emission features in the spectra of stars earlier than $\\sim$F0, such as $\\alpha$ CMa. Comparison of absolutely calibrated spectra of standard stars to corresponding spectra from the SWS reveals that these artifacts result from an underestimation of the strength of the CO and SiO molecular bands in the spectra of sources used as calibrators by the SWS. Although OLP 10.0 was intended to be the final data release, these findings have led to an additional release addressing this issue, OLP 10.1, which corrects the artifacts. ", + "introduction": "} The Short Wavelength Spectrometer (SWS) on the {\\em Infrared Space Observatory (ISO)} obtained approximately 1250 spectra covering the full 2.4--45~$\\mu$m wavelength range at moderate resolution. We are engaged in an ongoing project to classify these spectra \\citep{ksp01}, to reprocess them, and to present them in a publically available database on-line (Sloan et al. 2002, in preparation). SWS data are currently available in a partially processed form called Auto-Analysis Results (AARs). The AARs have a fairly complex format \\citep{lee01} that requires further processing before the data can be used for scientific analysis. Our processing method reduces these data to a single continuous 2.4 to 45~$\\mu$m spectrum. In assessing the quality of our new reprocessing algorithm, we compared the results for several infrared standards observed with the SWS to absolutely calibrated spectra from Cohen et al. (1992a, 1992b, 1995, 1996a, 1996b, 2001 (in preparation)). These standards are based on synthetic spectra of the A0 dwarf $\\alpha$ Lyr and the A1 dwarf $\\alpha$ CMa, which serve as the reference standards in the system. Cohen et al. (1992a) describe the details of their method that used high-quality ground-based and airborne photometry to normalize the synthetic spectra to measured astronomical fluxes. The synthetic spectra are based on the models of \\cite{kur79}, with updated opacities and metallicities \\citep{ck94}. Secondary standards are added to the system by observing their spectra in conjunction with that of the primaries so that atmospheric, telescopic, and instrumental transients can be removed. The spectra for the secondary standards are then obtained by dividing the observed secondary by the observed reference standard and multiplying by the assumed spectrum (the model) for the reference, degraded to match the spectral resolution of the instrument used: \\begin{eqnarray} S_{b,final}(\\lambda) = \\frac{S_{b,obs}(\\lambda)}{S_{a,obs}(\\lambda)} S_{a,assumed}(\\lambda), \\label{eqn.1} \\end{eqnarray} \\noindent where the subscript $a$ refers to the reference standard, and the subscript $b$ refers to the secondary. This is the standard method used by spectroscopists at ground-based telescopes to calibrate a program source by ratioing its spectrum to that of a standard star, preferably secured at an airmass matching that of the target. In effect, Equation \\ref{eqn.1} transfers the quality of the synthetic A star model from the reference to the new standard, whatever its spectral type. Cohen et al. (1992b, 1995, 1996a, 1996b, 2001) have applied this method to create absolutely calibrated composite spectra for 13 infrared standards, using spectra from ground-based telescopes, the Kuiper Airborne Observatory (KAO) and the {\\em Infrared Astronomical Satellite (IRAS)} \\citep[further details of the process can be found in ][]{cwb92a, cww92b}. \\cite{cwb96a} added $\\alpha^{1}$ Cen (G2 V) as an alternative reference standard for the southern hemisphere to give a total of 3 reference standards as well as the 13 secondaries. Most of the secondary standards are giants with spectral classes later than K0 and were chosen for their intrinsic brightness; sources later than approximately M3 are avoided due to the possibility of their variability \\citep{eg97}. The preferred infrared reference standard is Sirius ($\\alpha$ CMa), due to its brightness and its dust-free spectrum beyond 20~$\\mu$m. Figure 1 compares our derived SWS spectrum of this source to the Kurucz model presented by Cohen et al. (1992a). Deviations between the two occur in the vicinity of 4.5~$\\mu$m and again at 8.0~$\\mu$m. We propose an explanation for the origin of this discrepancy in this Letter and discuss the implications for the calibration of the SWS database and the impact on calibration of future infrared missions. ", + "conclusions": "We have identified spectral artifacts in the vicinity of the CO and SiO bands of late-type standard stars that appear in all OLP 10.0 spectra. These artifacts constitute a known systematic bias over a fairly large spectral range, and removing them is relatively straightforward. While OLP 10.0 was to have been the final version, an additional release, OLP 10.1, has been made to correct the database based on our findings. The SWS calibrations between 2 and 12~$\\mu$m were predominantly based on stellar atmospheric models of cool stars. These models have been used with great success in analysis of physical properties of cool stars \\citep[e.g.][]{dwe00}. Model spectra of A stars, which contain only atomic lines and no molecular bands, appear to be well founded. Synthetic spectra can also achieve higher spectral resolution than that afforded by the composites. However, the artifacts in the SWS database that mirror the dominant molecular absorption features in the 4--10~$\\mu$m range indicate that synthetic spectra of cool giants require further progress before they can be used for definitive spectral calibration. The Infrared Spectrograph (IRS) on the upcoming Space Infrared Telescope Facility (SIRTF) faces issues similar to those encountered by the SWS team, with the added complication that all of the commonly used spectral standards will be far too bright for use on the IRS. Until the commission of the Stratospheric Observatory for Infrared Astronomy (SOFIA), it will prove difficult to observationally calibrate standards faint enough for use by the IRS. Until this happens, the IRS will need to rely on models for spectral calibration. If models of later stars are required, improvements in the production of synthetic spectra will also be necessary." + }, + "0112/astro-ph0112392_arXiv.txt": { + "abstract": "We investigate the production of electron-positron pairs by inverse Compton scattered (ICS) photons above a pulsar polar cap (PC) and calculate surface heating by returning positrons. This paper is a continuation of our self-consistent treatment of acceleration, pair dynamics and electric field screening above pulsar PCs. We calculate the altitude of the inverse Compton pair formation fronts, the flux of returning positrons and present the heating efficiencies and X-ray luminosities. We revise pulsar death lines implying cessation of pair formation, and present them in surface magnetic field-period space. We find that virtually all known radio pulsars are capable of producing pairs by resonant and non-resonant ICS photons radiated by particles accelerated above the PC in a pure star-centered dipole field, so that our ICS pair death line coincides with empirical radio pulsar death. Our calculations show that ICS pairs are able to screen the accelerating electric field only for high PC surface temperatures and magnetic fields. We argue that such screening at ICS pair fronts occurs locally, slowing but not turning off acceleration of particles until screening can occur at a curvature radiation (CR) pair front at higher altitude. In the case where no screening occurs above the PC surface, we anticipate that the pulsar $\\gamma $-ray luminosity will be a substantial fraction of its spin-down luminosity. The X-ray luminosity resulting from PC heating by ICS pair fronts is significantly lower than the PC heating luminosity from CR pair fronts, which dominates for most pulsars. PC heating from ICS pair fronts is highest in millisecond pulsars, which cannot produce CR pairs, and may account for observed thermal X-ray components in the spectra of these old pulsars. ", + "introduction": "In the last several years, the basic model of particle acceleration above a pulsar polar cap (PC) has been undergoing significant revision. Sturrock (1971), Ruderman \\& Sutherland (1975) and Arons \\& Scharlemann (1979) originally proposed that particles are accelerated by an induced electric field, producing curvature radiation (CR) photons which create electron-positron pairs in the strong magnetic field. The pairs short out the electric field above a pair-formation front (PFF), self-limiting the acceleration. In the process of screening or shorting-out the electric field, some fraction of the positrons is accelerated back toward the PCs and heats the surface of the neutron star (NS), producing a potentially observable X-ray emission component (Arons 1981). Two recent developments have introduced important changes to this picture. First was the finding of Muslimov \\& Tsygan (1992) that the effect of inertial-frame dragging near the NS surface greatly increases the induced electric field $E_{\\parallel}$ above the PC in space-charge limited flow models (Arons \\& Scharlemann 1979). The second was the realization that inverse-Compton radiation of the primary particles can produce pairs which could potentially screen the electric field (Zhang et al. 1997, Harding \\& Muslimov 1998; hereafter HM98). We have investigated the electric field screening and PC heating in the revised space-charge limited flow (SCLF) model. In the first paper (Harding \\& Muslimov 2001; Paper I), we have presented our results for screening and PC heating by CR PFFs, as in Arons (1981), but using the frame-dragging electric field of Muslimov \\& Tsygan (1992). Paper I outlined a self-consistent calculation of the PFF height, returning positron flux and the screening scale length, where the $E_{\\parallel}$ and the primary flux are adjusted for the change in charge density caused by the returning positrons. A main assumption of this calculation was that the screening scale is small compared to the height of the PFF, the location where the first pairs are produced. This turns out to be generally true for screening by CR-produced pairs because the pair cascade multiplicity grows rapidly over small distances due to the strong dependence of CR photon energy on particle energy. Thus, the existence of a CR PFF always results in PC heating and full screening of $E_{\\parallel}$. We found that the PC heating luminosity, as a fraction of the spin-down luminosity, increases with pulsar characteristic age, $\\tau = P/2\\dot P$, and should be detectable for pulsars with $\\tau \\gsim 10^6$ yr. The most significant heating occurs for pulsars near the death line for CR pair production, which is $\\tau \\sim 10^7$ yr for normal period pulsars and $\\tau \\sim 10^8$ yr for millisecond pulsars. Our predicted X-ray luminosity due to PC heating is about a factor of ten higher than the X-ray luminosity predicted by Arons (1981), due to the increase in accelerating voltage drop resulting from the inclusion of inertial-frame dragging effects. We also predicted that older pulsars should have higher PC surface temperatures from heating. In this second paper (Paper II), we present results of our investigation of electric field screening and PC heating by inverse-Compton scattering (ICS) radiation PFFs. We investigate PC heating by ICS produced pairs in all pulsars, including those which produce CR pairs. For the older pulsars that do not produce pairs through CR, PC heating by positrons from ICS cascades is especially important. Our treatment of the ICS PFFs follows closely that of HM98, where cyclotron-resonant and non-resonant scattering are considered as separate radiation components. HM98 found that the ICS PFFs are located much closer to the NS surface than are the CR PFFs, because primaries with Lorentz factor of only $\\sim 10^4 - 10^5$ can produce pairs via ICS whereas Lorentz factors of at least $\\sim 10^7$ are required for production of pairs via CR. They also found that the ICS PFFs of positrons returning from the upper PFF occur at a significant distance above the surface, so it is possible that ICS pairs fronts are unstable if screening of $E_{\\parallel}$ at the lower PFF occurs. However, HM98 did not determine whether the ICS pairs were capable of screening either at the upper or lower PFFs. This paper will attempt to answer those questions with a calculation similar to that of Paper I. In Section II, we describe the acceleration model for this calculation. We use a solution for $E_{\\parallel}$ from Poisson's equation which imposes an upper boundary condition requiring that $E_{\\parallel} = 0$ at the location of screening (as in HM98), which differs from that of Paper I which used a solution with no upper boundary. The solution with an upper boundary is more appropriate for screening at ICS PFFs which form close to the surface, where the screening scale is generally comparable to the PFF height. This solution is required here because pairs from ICS photons may screen $E_{\\parallel}$ close to the NS surface. In this case the value of $E_{\\parallel}$ is suppressed due to proximity of the screening to the surface. Section II also presents a calculation of ICS PFF height and the location of a new ICS pair death line. We find that virtually all known pulsars can produce pairs by either resonant ICS (RICS) or non-resonant ICS (NRICS) in a pure dipole field. Hibschmann \\& Arons (2001) (hereafter HA01) have recently taken a different approach to determining pulsar death-lines including ICS produced pairs. They define the PFF as the location where the pair multiplicity achieves that required for complete screening of the $E_{\\parallel}$, whereas we define the PFF as the location where pair production begins (i.e. where the first pairs are produced). In the case of CR, this distinction is minor since the screening scale is small compared to the PFF height, but in the case of ICS the distinction is very important since the screening scale is comparable to the PFF height. HA01 also assume that the pair multiplicity required for screening is determined by the difference between the actual charge and the Goldreich-Julian charge at a distance from the NS surface that is roughly equal to the PC radius, whereas in our calculation the charge density required for screening is determined by the charge deficit at the location of creation of pairs which can be much smaller. With their more restrictive definition of the PFF, their death lines differ significantly from ours. In Section III, we give self-consistent solutions for the fraction of positrons returning to the PC and for the screening scale. We also explore the question of whether positrons returning from the upper PFF can screen $E_{\\parallel}$ near the lower PFF close to the NS surface. The subject of pulsar pair death lines, i.e. which NSs are capable of pair production, will be discussed in Section IV. In Section V, we present our calculations of PC heating luminosity due to returning positrons, giving both numerical and analytic estimates. Summary and conclusions, as well as a comparison of results from both Papers I and II will be given in Section VI. ", + "conclusions": "We have explored production of electron-positron pairs by photons produced through ICS of thermal X-rays by accelerated electrons above a pulsar PC. Since the accelerating primary electrons can produce pairs from ICS photons at much lower energies than are required to produce pairs from CR photons, it is very important to investigate the consequences of ICS pair fronts for $E_{\\parallel}$ screening and for PC heating. We have defined ``pair death\" lines in $B_0 - P$ space as the boundary of pair production for pulsars, having dipole magnetic fields of surface strength $6.4 \\times 10^{19}(\\dot P P)^{1/2}$ G. Operationally, the existence of a pair front is determined by a finite solution to equation (\\ref{S0}) for $S_0$, the altitude of the onset of pair creation. Although we are able to give analytic formulae for the altitude of the PFFs for different radiation processes, the location of the pair death lines must be determined numerically. The existence of a pair front requires much less than one pair per primary electron (but still many pairs from the whole PC beam), since the very first pairs are created in the declining high-energy tail of the radiated spectrum, and also much less than the number of pairs per primary required for screening of $E_{\\parallel}$. We find that virtually all known radio pulsars are capable of producing pair fronts with ICS photons. A smaller number, less than half, are able to also produce pairs via CR. If the acceleration model we use is correct and pulsars have dipole fields, then this result implies that relatively few pairs are required for coherent radio emission. Self-consistent calculations show that ICS pair fronts produce lower fluxes of returning positrons and lower PC heating luminosities than CR pair fronts. This is due to the higher efficiency of the ICS process in producing pairs at lower altitudes where both the charge deficit required to screen (and thus the returning positron flux) and the accelerating voltage drop are smaller. For pulsars with surface magnetic fields in the ``normal\" range of $10^{11} - 10^{13}$ G, ICS heating luminosities are several orders of magnitude lower than CR heating luminosities. However, for ms pulsars having surface fields in the range $10^8 - 10^{10}$ G, production of any pairs requires such high photon energies that ICS pair fronts occur at higher altitudes, where acceleration voltage drops are high enough to produce significantly more PC heating. Since most ms pulsars cannot produce pairs through CR, ICS pair fronts provide the only means of external PC heating. We find that for surface temperatures $T \\gsim 10^6$ K, ICS heating luminosities are in the range of detection. We find that ICS pairs are able to screen the {\\it local} $E_{\\parallel}$ in some pulsars having a high enough PC temperature, but that this local screening will not produce a complete screening of the accelerating field at all altitudes and thus will not stop the acceleration of the primary beam. This is because the number of ICS pairs grows slowly, on a scale length comparable to the altitude of the PFF, and then declines while the charge deficit which maintains $E_{\\parallel}$ continues to increase with altitude. Even if the ICS pair production is vigorous enough to achieve local screening of $E_{\\parallel}$, it eventually cannot produce the charge density (from returning positrons) to keep up with the increasing charge deficit (produced by the combination of flaring field lines and inertial frame-dragging). The primary particles may slow their acceleration briefly, due to the local screening, but will resume acceleration once the ICS pair production declines. At higher altitudes, those pulsars to the left of the CR pair death line will reach the Lorentz factors ($\\gamma \\sim 2 \\times 10^7$) required to produce CR pairs. In contrast to ICS pair fronts, the growth of pairs above the CR PFF is rapid and robust due to the sensitivity of CR photon energy and emission rate on particle Lorentz factor, producing complete screening of $E_{\\parallel}$ in a very short distance (cf. Paper I). In pulsars to the right of the CR pair death line, there is not complete screening and the primary particles will continue accelerating to high altitude with their Lorentz factor being possibly limited by CR reaction. We have also investigated the proposal by HM98 that pairs produced as the returning positrons are accelerated toward the NS may be able to screen the $E_{\\parallel}$ above the surface. Using our calculated values of the returning positron flux in cases where local screening has been achieved at the ICS pair fronts, we find that in most cases screening does not occur at a significant distance above the NS surface to cause disruption of a steady state or formation of pair fronts. In the cases where screening does occur near PFFs significantly above the surface, the pulsars are near the ICS screening boundary. The resulting instability would then not move the start of the acceleration to higher altitudes, as HM98 had envisioned, but would probably weaken or disrupt the screening at the upper ICS pair front, resulting in a decrease in returning positron flux which would weaken the screening at the lower PFF, etc. PC heating by CR pair fronts will dominate for pulsars to the left of the CR pair death line, while ICS pair fronts will supply the PC heating for pulsars to the right of the CR pair death line. While we have given analytic expressions for the fraction of returning positrons and PC heating luminosities from ICS pair fronts, these are only good estimates above the ICS screening boundaries of Figure 4. Below the ICS screening boundaries, where local screening is not achieved for that PC temperature, the numerical values of returning positron fraction and heating luminosity fall well below the analytic estimates. This will be true for ms pulsars, nearly all of which are below the ICS screening boundary for a PC temperature of even $T = 10^7$ K. Our results are dependent on a number of assumptions inherent in our calculations. First, we have assumed that the pattern of thermal X-ray emission is from a heated PC and is a pure, isotropic blackbody. According to recent calculations (Zavlin et al. 1995) of radiation transfer in magnetized NS atmospheres, the thermal emission is not pure blackbody or isotropic, but a somewhat cooler blackbody consisting of pencil and fan beam components. Both effects of full surface (cooler) emission and pencil beaming would tend to decrease ICS radiation and thus ICS pair production efficiency. However, a large fan beam component would tend to increase ICS efficiency. Second, we have used a hybrid scheme to describe the ICS radiation spectrum in which the RICS has been treated as classical magnetic Thompson scattering and NR (or continuum) ICS has been treated as relativistic but non-magnetic. In reality, both RICS and NRICS should be treated as a single process with one cross section. While the magnetic QED scattering cross section has been studied for some time (e.g. Herold 1979, Daugherty \\& Harding 1986), simple expressions in limited cases are only beginning to become available (e.g. Gonthier et al. 2000). Our present treatment is probably accurate for magnetic fields $B \\lsim 0.2 B_{\\rm cr}$ which includes most of the radio pulsars. The location of a pulsar relative to the pair death lines may be important not only to its radio and thermal X-ray emission characteristics, but also to its high-energy emission properties. As we have argued in this paper, ICS pair fronts will not limit the acceleration voltage drop in pulsars but that acceleration will continue until it is limited by a CR pair front. The voltage drop at the CR pair front (together with the size of the PC current) is therefore expected to determine the high-energy emission luminosity. In Section 4, we have noted that the acceleration voltage drop of pulsars that produce CR pair fronts is remarkably insensitive to pulsar parameters, leading to the prediction that high-energy luminosity, $L_{\\sc HE}$, should be simply proportional to PC current (which is proportional to $\\dot E_{\\rm rot}^{1/2}$), which seems to be borne out by observations (e.g. Thompson 2000). However, pulsars that do not produce CR pair fronts do not have such a limit on acceleration voltage drop and should depart from the $L_{\\sc HE} \\propto \\dot E_{\\rm rot}{1/2}$ dependence, and approach a $L_{\\sc HE} \\propto \\dot E_{\\rm rot}$ dependence. Indeed, such a departure must occur if they are not to exceed 100\\% efficiency in converting rotational energy loss to high-energy emission. We predict that this change in $L_{\\sc HE}$ dependence should occur along the CR pair death line. None of the pulsars which have detected $\\gamma$-ray emission are to the right of the CR pair death line although some, such as Geminga, are close. The Large Area Gamma-Ray Telescope (GLAST) will have the sensitivity to detect $\\gamma$-ray emission from significant numbers of radio pulsars beyond the CR pair death line and so should be able to test this prediction." + }, + "0112/astro-ph0112501_arXiv.txt": { + "abstract": "I highlight the remarkable advances in the past few years in CMB research on total primary anisotropies, in determining the power spectrum, deriving cosmological parameters from it, and more generally lending credence to the basic inflation-based paradigm for cosmic structure formation, with a flat geometry, substantial dark matter and dark energy, baryonic density in good accord with that from nucleosynthesis, and a nearly scale invariant initial fluctuation spectrum. Some parameters are nearly degenerate with others and CMB polarization and many non-CMB probes are needed to determine them, even within the paradigm. Such probes and their tools were the theme of the TAW8 meeting: our grand future of CMB polarization, with AMiBA, ACBAR, B2K2, CBI, COMPASS, CUPMAP, DASI, MAP, MAXIPOL, PIQUE, Planck, POLAR, Polatron, QUEST, Sport/BaRSport, and of Sunyaev-Zeldovich experiments, also using an array of platforms and detectors, \\eg AMiBA, AMI (Ryle+), CBI, CARMA (OVROmm+BIMA), MINT, SZA, BOLOCAM+CSO, LMT, ACT. The SZ probe will be informed and augmented by new ambitious attacks on other cluster-system observables discussed at TAW8: X-ray, optical, weak lensing. Interpreting the mix is complicated by such issues as entropy injection, inhomogeneity, non-sphericity, non-equilibrium, and these effects must be sorted out for the cluster system to contribute to ``high precision cosmology'', especially the quintessential physics of the dark energy that adds further mystery to a dark matter dominated Universe. We will have to address ``Is it cluster evolution or is it cosmology?''. The answer will be both, but we can be optimistic that, with the huge data influx, computational power increase, and talented people joining the adventure, we can handle both observationally, theoretically and phenomenologically. ", + "introduction": "\\subsection{The Beginning of the End or the End of the Beginning?} In April 2001, just predating TAW8\\footnote{This paper blends an introductory primary CMB talk with my conference summary. Only a few CMB references are given, organized by date (April'99, April'00, April'01). The perpetrators of the advances mentioned in this summary, and associated references, can be found elsewhere in these proceedings.}, the Boomerang and DASI teams independently unveiled remarkably similar power spectra of the {\\it primary} anisotropies of the CMB, those which can be calculated using linear perturbation theory (Fig.~\\ref{fig:CLopt}). The analysis of cosmological parameters was in accord with indications from Large Scale Structure (LSS), Supernova (SN1), and a variety of other observations, pointing towards everyone's neo-standard model at this meeting, $\\Lambda$CDM. Typical $\\Lambda$CDM parameters are taken to be: $\\Omega_{tot}=1$; $\\Omega_\\Lambda \\approx 0.7$; Hubble parameter $h \\approx 0.7$ from the Hubble key project; $\\Omega_m \\sim 0.3$, including $\\sim 0.04$ in baryons, the rest in cold dark matter; $n_s$=1 as the slope of the initial density power spectrum, the scale-invariant Harrison-Zeldovich-Peebles value; overall mass density power normalized to have $\\sigma_8 \\sim 0.9$, with $\\sigma_8$ the {\\it rms} (linear) density fluctuation level on a cluster-scale ($8\\hmpc$). The baryon density choice $\\Omega_{b}h^2 \\approx 0.02$ is the Big Bang Nucleosynthesis result calibrated with the deuterium abundance estimated from absorption lines in QSO spectra. From the early 80s onward, CMB observations were used with LSS information, as embodied in angular and redshift galaxy surveys, cluster and other rare event abundances, cluster clustering, and velocity flows to constrain the cosmological parameters defining the space. Even in the days of CMB upper limits predating the COBE/FIRS/SP91/Tenerife and subsequent detections, the CMB was a powerful constrainer. When the COBE detection was combined with LSS, a great collapse occurred in parameter space, which was further constricted by detections on intermediate angular scales throughout the 90s, and which Boomerang, DASI and Maxima have now turned into bulls-eye determinations on some key parameters (Fig.~\\ref{fig:cmbLSS2sig}), focussing even more than in the April'00 release. \\begin{figure} \\plotone{fig1optTaiwan12_09.eps} \\vspace{-5pt} \\caption{The optimally-combined power spectrum ${\\cal C}_\\ell$ grouped in bandpowers using all current data (circles, joined by a light line) is contrasted with that for Boomerang-LDB (squares), DASI (crosses) and DMR (point at low $\\ell$). ``pre'' denotes TOCO, Boom-97 and 19 other experiments predating April'99. This heterogeneous ``prior CMB'' mix is quite consistent with what Boomerang, DASI and Maxima show, with much larger errors. CBI2 denotes the two published CBI points, only a small fraction of the total CBI data. A Boomerang best-fit model using the weakH+LSS+flatU prior is also shown. In spite of the 10\\% calibration and 13\\% beam uncertainties for Boomerang, little adjustment of its median values was required by the other data. A caveat: DASI's fields overlap about 5\\% of the Boomerang area, so there is correlation between Boomerang and DASI. This is not taken into account here, but the consistency in the overlap regions are currently being explored. The optimal ${\\cal C}_\\ell$ without DASI included looks similar to the one shown, as might be expected given the consistency of the two power spectra (and also the similarity in derived cosmic parameters). } \\label{fig:CLopt} \\end{figure} \\begin{figure} \\plotone{2sigBPdasimshtfnoqso_precmb.eps} \\vspace{-10pt} \\caption{2-$\\sigma$ likelihood contours for the dark matter density $\\omega_c =\\Omega_{cdm}{\\rm h}^2$ and $\\{\\Omega_k,\\Omega_\\Lambda,n_s,\\omega_b\\}$ for the LSS+weakH prior, and the following CMB experimental combinations: DMR (short-dash); the ``April'99\"+DMR data (short-dash long-dash); TOCO + (April'99+DMR) data (dot short-dash); ``prior-CMB\" = Boom-97 + (TOCO+April'99+DMR) data (dot long-dash); Boomerang + DASI + Maxima-1 + ``prior-CMB\" data (heavy solid, all-CMB). These $2\\sigma$ lines tend to go from outside to inside as more CMB experiments are added. The smallest 2-$\\sigma$ region (dotted and interior) shows SN1+LSS+weakH+all-CMB, when SNI data is added. For the $\\Omega_\\Lambda$, $n_s$ and $\\omega_b$ plots, the flatU prior, $\\Omega_{tot}$=1, has also been assumed, but the values do not change that much if $\\Omega_{tot}$ floats. The main movement from Apr'00 to Apr'01 was that $\\omega_c$ localized more around 0.13 in all panels, and the $\\omega_b$ contour in the lower right panel migrated downward a bit to be in its current good agreement with Big Bang Nucleosynthesis. } \\label{fig:cmbLSS2sig} \\end{figure} It appears from Fig.~\\ref{fig:CLopt} that multiple peaks and dips in the CMB have been found -- a dominant first peak, a less prominent second one, and a hint of a third one, with interleaving dips (April'01). These are even in roughly the right location of a long-standing prediction of adiabatic inflation-based models with little mean curvature. The physics of the ${\\cal C}_\\ell$ peak structure is based on acoustic oscillations and velocity flows as the photon-baryon fluid viscously passed from tight coupling to free-streaming at photon decoupling (redshift $z_{dec} \\sim 1100$, about 0.4 Myr after the ``Big Bang''), generating the ``damping tail'' evident in the ``best-fit'' theoretical model shown in Fig.~\\ref{fig:CLopt}. The maps from which the ${\\cal C}_\\ell$ bandpowers are derived are largely noise-free images of soundwave patterns seen through the photon decoupling \"surface\" of width $\\sim 10 \\hmpc$ that defines the thick-to-thin transition. This isquite a bit smaller than the comoving \"sound crossing distance\" at decoupling, $\\sim 100 \\hmpc$ (\\ie $\\sim 100 \\kpc$ physical), below which density oscillations and velocity flows can be observed. After, photons freely-streamed along geodesics to us, mapping (through the angular diameter distance relation) the post-decoupling spatial structures in the temperature to the angular patterns we observe now. The free-streaming along our (linearly perturbed) past light cone leaves the pattern largely unaffected, except that temporal evolution in the gravitational potential wells as the photons propagate through them leaves a further $\\Delta T$ imprint, called the integrated Sachs-Wolfe effect. Of course there are a number of other signals that are also present in the maps, so how can we be confident that Fig.~\\ref{fig:CLopt} really offers a glimpse of fluctuation power at $z_{dec}$? Known contaminating signals include the Galactic foregrounds of bremsstrahlung, synchrotron and dust emission, extragalactic radio and infrared sources. As well, {\\it secondary} anisotropies associated with post-decoupling nonlinear effects are also present, These include weak-lensing by intervening mass, Thompson-scattering by the nonlinear flowing gas once it became \"reionized\" at $z \\sim 10-20$, the thermal and kinematic SZ effects, and the red-shifted emission from dusty galaxies. All secondary and foreground sources leave non-Gaussian imprints on the CMB sky, and all but the kinematic SZ effect have different spectral signatures to aid in signal separation. For some experiments (DASI, CBI), it has been crucial to remove sources, for others like Boomerang, relatively contamination-free channels and regions can be found. We have been lucky that many of these signals are subdominant at the angular scales we are probing in Fig.~\\ref{fig:CLopt}. As precision improves, signal separation will loom large. Because of the CMB+LSS success, we did not see at TAW8 as many of the usual comparison cosmologies as we used to at such meetings, the open oCDM, the hot/cold hybrid HCDM, $\\tau$CDM, the tilted tCDM, the cluster-normalized old-standard sCDM. Nor were cosmic defect models in evidence. Though many of the $x$CDM's may have fallen away, we now see $Q$CDM appearing on the stage, with $Q$ an ultra-low mass scalar field, often called quintessence, that dominates at late times. Thus $\\Omega_Q$ replaces $\\Omega_\\Lambda$, and an effective $Q$-dynamics is cast (though none too well) in terms of a mean pressure-to-density ratio $w_Q = \\bar{p}_Q/\\bar{\\rho}_Q$, an effective equation of state (EOS). Well not so effective, since $Q$ is a spatially-dependent field, or may be. In spite of a huge number of quintessential papers, $Q$ would better stand for question mark. For $\\Lambda$, $w_Q=-1$, but $w_Q < - 1/3$ would get our patch of the Universe into acceleration, apparently with no new comoving space to be revealed. If there really is a $\\Lambda$CDM/$Q$CDM concordance, then apart from the wide grins of the ``often in error, never in doubt'' cosmologists, hubrous abounding, we may also hear theorists' lament: Where are the anomalies for wild and fun theorizing? Between our state now, with its large-ish error bars and the never-ending worry about the systematic rather than the statistical, and with the exquisite data from a vast array of experiments coming down the pipe, there is still much room for a cosmic surprise. Perhaps the greatest of all will be if the models of the 80's do in fact describe how all of the structure formed in the Universe, albeit with a mysterious dark energy accelerating us. Even if $\\Lambda$CDM, theorists are still at play, though not so much at TAW8 which was concretely directed to the empirical. Just look to the dark energy, the struggles to tie the latest inflation our observable patch of the Universe now seems to be caught in with the early inflation needed to ``smooth the universe'' and solve causality problems, and, incidentally, to generate quantum noise from which all observed cosmic structure originated. Look to the dialogues between those of the M-theory brane worlds and the physical cosmologists, reigniting the early universe connection that we were in danger of losing -- what with the (clustering) dark matter being supposed cold for so long and with inflation being generic but tunable to meet all demands (though not without highly baroque additions). \\subsection{Broad Truths from the CMB+LSS} Most amazing about Fig.~\\ref{fig:CLopt}, COBE's FIRAS experiment, the accumulating LSS information, now coming in a torrent with 2dF, Sloan and higher redshift surveys, \\etc, is that the paradigm appears to hold: a hot Big Bang, with an almost perfect $T_{\\gamma *} =2.725 \\pm 0.001$K blackbody spectrum that must have come to us from beyond the most distance SZ cluster, $z \\gta 1$. That ${\\cal C}_\\ell$ is significantly positive at $\\ell \\sim 1000$ argues against a large $\\exp[-2\\tau_C]$ damping multiplier, where $\\tau_C \\sim 0.1 (\\omega_b/0.02) (\\omega_m/0.15)^{-1/2} ((1+z_{reh})/15)^{3/2}$ is the Thompson optical depth to the epoch $z_{reh}$ of reheating. Thus, though much pregalactic energy injection at $z \\sim 200$ is still possible, it does not look like it. The FIRAS limit of $4\\bar{y} < 10^{-4}$ on fractional energy input into the CMB from the lack of a Compton cooling spectral $y$-distortion further implies no large entropy injection could have occurred at lower $z$ into the gas, strongly limiting the role explosions can have had in LSS development. The beautful direct connection of the small $\\Delta T$ fluctuations to the density amplitudes now -- on the same spatial scales -- strongly support the gravitational instability picture of structure formation. That it forms hierarchically, from small to big, is of course obvious from LSS observations at various redshifts, but $n_s \\sim 1 $ from the CMB adds further positive support. The primary CMB fluctuations are quite Gaussian, according to COBE, Maxima, and now Boomerang analyses. A non-Gaussian component, possibly subdominant, of the primordial fluctuations can still work, but it is encouraging for inflationists where Gaussian statistics are the natural (but not only) outcome. Cosmic defect and cosmic string models of structure formation are more challenged by the peaks and dips of ${\\cal C}_\\ell$, which are very difficult to get, than by the Gaussianity. We know the gravitational instability of a hierarchical Gaussian random density field leads naturally to the cosmic web interconnections and the prevalence of superclustering that we seem to find observationally at low and high redshifts --- a framework for thinking about the cluster/group system that was a main theme of TAW8. The web consists of massive clusters with overdensities $\\delta \\gta 100$, filaments with $\\delta \\sim 5-10$, which bridge massive clusters, groups which bead the bridges, membranes with $\\delta \\sim 2$ which join the filaments, and the voids with $\\delta < 0$ dominating the space but not the mass. This picture is of course borne out by all the large simulations reported at TAW8, sizes ranging from $128^3$ for a ``Schrodinger equation'' cosmological calculation to $256^3$ and even $512^3$ for hydro and $N$-body, and to $1000^3$ for $N$-body. Just a decade ago, a $128^3$ $N$-body was a tour de force. We also heard much about semi-analytic methods in many different guises that fit into this web picture, the halo model and the peak-patch model with clustering included in both, and of course many variants of Press-Schechter-ism. ", + "conclusions": "" + }, + "0112/astro-ph0112098_arXiv.txt": { + "abstract": "We investigate the structure and the stabilities of a protoplanetary disk, which is heated by viscous process in itself and by its central star. The disk is set to rotate with the Keplerian velocity and has the surface density distribution of the minimum mass solar nebula. We assume the vertical hydrostatic equilibrium and the radiative equilibrium at each point, and solve the two-dimensional radiative transfer equation by means of the Short Characteristic method in the spherical coordinate in order to determine the disk structure. Our calculation shows that at the outer region of the disk with a distance from the central star of $x>1$AU the radiative heating from the inner disk dominates the viscous heating even near the midplane. It is because of the high temperature distribution in the optically thin surface layer and the relatively high disk height ($z_{\\infty}\\sim 0.7x$ at $x\\sim 1$AU) as a consequence of the irradiation from the inner hot region of the disk. In addition, we examine the convective and the magnetorotational instabilities of the disk. As a result, the whole disk is convectively stable since the dusty region is not heated by the viscous dissipation from the midplane but by the radial radiative heating. On the other hand, almost all the disk is magnetorotationally unstable except for the region near the equatorial plane of 2AU $1$AU, the radiative heating from the inner disk dominates the viscous heating because the temperature at the surface region is high and the disk is geometrically thick, owing to the irradiation from the inner hot region of the disk. Making use of the resultant density and temperature profiles, we have examined the convective and the magnetorotational instabilities of the disk, which are expected to induce the turbulent motion, and then, bring about the angular momentum and mass transfer, and/or the energy transport. As a result, we have found that: 3. The whole disk is convectively stable since the dusty region is not heated by the viscous dissipation from the midplane but the radiative heating from the inner disk. 4. Almost all of the disk is magnetorotationally unstable as the ionization degree is high enough that the stabilizations by ambipolar diffusion and ohmic dissipation are negligible, except for the region near the midplane of 2AU $ -15$) appears to be higher (perhaps much higher) than one fifth but it remains very uncertain. The fraction that are overluminous ($M_B < -20$) is lower (probably much lower) than 0.01. The absolute-magnitude distributions for each supernova type, restricted to events within 1 Gpc, are compared. Although these distributions are affected by observational bias in favor of the more luminous events, they are useful for comparative studies. We find mean absolute blue magnitudes (for $H_0=60$) of $-19.46$ for normal Type~Ia supernovae (SNe~Ia), $-18.04$ for SNe~Ibc, $-17.61$ and $-20.26$ for normal and bright SNe~Ibc considered separately, $-18.03$ for SNe~II-L, $-17.56$ and $-19.27$ for normal and bright SNe~II-L considered separately, $-17.00$ for SNe~II-P, and $-19.15$ for SNe~IIn. ", + "introduction": "The absolute-magnitude distributions of the various supernova (SN) types provide vital information for determining SN rates, for advancing our knowledge of the stellar progenitors and their explosion mechanisms, and for planning future ground- and space-based SN searches. More than a decade ago, \\citet[MB90;][]{mb90} used the Asiago Supernova Catalog \\citep[the ASC;][]{barbon89} to carry out a comparative study of absolute-magnitude distributions. At the end of 1989, the ASC listed only 687 SNe, and many of those were of unknown type. By June, 2001 (updates are available at \\verb+http://merlino.pd.astro.it/~supern/+), the number of events in the ASC had increased to 1910. Almost all of the SNe of the 1990's had been assigned types, and some of the pre-1990 data had been improved. We therefore decided that an updated study of the absolute-magnitude distributions would be timely. It should be acknowledged that the characteristic uncertainty of the apparent magnitudes listed in the ASC, perhaps 0.2 or 0.3 mag, is not negligible for this purpose. Thanks to the tremendous interest in using SNe~Ia as distance indicators for cosmology, studies of the SN~Ia absolute-magnitude distribution now can be based on carefully selected samples of events for which both the apparent magnitudes and the relative distances are known to high accuracy \\citep[e.g.][]{hamuy96}. The absolute-magnitude scatter in these samples can be further reduced in two ways. First, the host galaxy extinction can be taken into account for each SN using the SN's $B-V$ color. Next, the fact that some SNe~Ia are intrinsically dimmer than others can be accounted for by using a correction to the light-curve width such as the $\\Delta m_{15}$ parameter \\citep{phillips99} or the method of multicolor light-curve shapes \\citep[MLCS,][]{riess96}. With these corrections the dispersion in $M_{B}$ for SNe~Ia can be reduced to 0.11 \\citep{phillips99}, making them extremely useful as distance indicators for cosmology. Our present study, which has little to add to our knowledge of the SN~Ia absolute-magnitude distribution, is directed more at the absolute-magnitude distributions of the other SN types. As far as the other SN types are concerned, either the number of events for which accurate peak apparent magnitudes have been reported remains small (SNe~Ib, Ic, II-L, IIn) or the intrinsic dispersion in the peak absolute magnitude is large (SNe~II-P). For these types a study such as this one, based on all available data, can be useful. For comparison with the other types, SNe~Ia are included as well. ", + "conclusions": "At least 7 of 31 SNe in our Galaxy and in galaxies within 10 Mpc appear to have been subluminous ($M_B \\geq -15$). Considering that even in this sample there is an observational bias against them, it appears that more (perhaps much more) than one fifth of all SNe are subluminous, but because the number of such events seen so far is small, this fraction remains very uncertain. Only 20 of 297 extragalactic maximum-light SNe appear to be overluminous ($M_{B}<-20$). Considering the strong observational bias in favor of them, and that observational errors produce spuriously overluminous events, it is safe to conclude that the fraction of all SNe that are overluminous must be much lower than 0.01. The results of our comparative study of the absolute-magnitude distributions of SNe in the distance-limited sample ($\\mu<40$) are summarized in Table~1. The main differences between our results and those of MB90 are that (1) it has become clear that overluminous SNe ($M_B \\leq -20$) exist; (2) the absolute-magnitude dispersion of SNe~Ibc has increased due to the discovery of some rather luminous events; and (3) we present results for SNe~IIn, which on average are the most luminous type of core--collapse SNe. Even though the SN discovery rate increased dramatically in the 1990's, the numbers of subluminous events, overluminous events, SNe~II-L, SNe~Ibc, and SNe~IIn for which apparent magnitudes at maximum light are available are still small --- this is especially so if SNe~II-L and Ibc should be divided into two luminosity groups. Except for SNe~Ia and II-P, the study of SN absolute-magnitude distributions remains data starved. Systematic programs of discovery and observation of events in the Hubble flow, such as the Nearby Supernova Factory \\citep{alder01} underway at the Lawrence Berkeley National Laboratory, should substantially improve the situation during the coming years. \\clearpage" + }, + "0112/astro-ph0112267_arXiv.txt": { + "abstract": "In recent years several Dwarf Novae (DN) systems have been observed in quiescence, when the accretion rate is low and the WD photosphere can be directly detected. The WDs are observed to cool after the DN outburst from high effective temperatures to lower effective temperatures ($T_{\\rm eff}$) thought to be indicative of the thermal state of the deep interior of the WD. Sion has argued that the most likely energy source for this quiescent luminosity is the gravitational compression of the WD interior, which rejuvenates an otherwise cold WD into a much hotter state. We are undertaking a theoretical study of the compressional heating of WD's, extending down to the very low time averaged accretion rates, $\\langle \\dot M \\rangle\\sim 10^{-11}M_\\odot \\ {\\rm yr^{-1}}$, applicable to the post-turnaround CV's (the ``TOADS''). Nuclear burning is unstable at these $\\langle \\dot M \\rangle$'s, so we have incorporated the recurrent heating and cooling of the WD throughout the classical novae limit cycle. In addition to self-consistently finding the range of $T_{\\rm eff}$ as a function of $\\langle \\dot M \\rangle$ during the cycle, we also self-consistently find the ignition masses. Comparing these theoretical masses to the observed ejected masses will tell us whether the WD mass in CV's is secularly increasing or decreasing. We close by comparing our results to the accumulated observations of quiescent DN and making predictions for the colors of low $\\langle \\dot M \\rangle$ CV's in quiescence that are applicable to searches for faint CVs in the field and galactic globular clusters. ", + "introduction": "Dwarf Novae (DN) systems contain a white dwarf (WD) accreting matter at time-averaged rates $\\langle \\dot M \\rangle<10^{-9}M_\\odot \\ {\\rm yr}^{-1}$ from a low-mass ($<0.5M_\\odot$ typically) stellar companion. At these $\\langle \\dot M \\rangle$'s, the accretion disk is subject to a thermal instability which causes it to rapidly transfer matter onto the WD (at $\\dot M \\gg \\langle \\dot M \\rangle$) for a few days to a week once every month to year. The orbital periods of these binaries are usually less than 2 hours (below the period gap), but there are also DN above the period gap, $>$ 3 hours (see Shafter 1992). The $\\dot M$ onto the WD can be low enough between outbursts that the optical/UV emission is dominated by the internal luminosity of the WD, not the accretion disk. Recent HST/STIS spectroscopy has spectrally resolved the WD's contribution to the quiescent light and found effective temperatures $T_{\\rm eff}\\sim 10,000-40,000 \\ {\\rm K}$ (see our Figure \\ref{fig:lars} and Sion 1999). The measured internal WD luminosity is larger than expected from an isolated WD of similar age ($\\approx$ Gyr), indicating that it has been heated by accretion (Sion 1985). Compressional heating (i.e. internal gravitational energy release) appears to be the main driver for this re-heating (Sion 1995). Sion's (1995) early estimate for internal gravitational energy release within the WD (of mass $M$ and radius $R$) was $L\\approx 0.15 GM\\langle \\dot M \\rangle/R$. However, we show in \\S 2 that most energy is released in the accreted outer envelope, giving $L\\approx 3kT_c \\langle \\dot M \\rangle/\\mu m_p$, where $\\mu\\approx 0.6$ is the mean molecular weight of the accreted material, $m_p$ is the baryon mass, and $k$ is Boltzmann's constant. The theoretical challenge that we address in \\S 3 is how to calculate the WD core temperature, $T_c$, as a function of $\\langle \\dot M \\rangle$, and thus find $T_{\\rm eff}$. Because of the unstable nuclear burning and resulting classical novae cycle, the envelope mass changes with time. This allows the core to cool at low accumulated masses and be heated prior to unstable ignition. We use nova ignition to determine the maximum mass of the overlying freshly accreted shell, and find the steady-state (i.e. cooling equals heating throughout the classical novae cycle) deep interior temperature of the accreting WD, $T_c$, as a function of $\\langle \\dot M \\rangle$ and WD mass. In \\S 4, we compare our theoretical work to STIS observations and infer $\\langle \\dot M \\rangle$ on the timescale of $10^6$ years, critical to constraining CV evolutionary models. Figure \\ref{fig:lars} shows that DN above the period gap are hotter than those below the gap, and have $\\langle \\dot M \\rangle$'s consistent with that expected from traditional CV evolution (e.g. Howell et al. 2001), even those that involve some ``hibernation'' (Shara et al. 1986; Kolb et al. 2001). The result is more surprising if the much weaker magnetic braking laws of Andronov et. al. (2001) are correct. We also predict the minimum light ($M_V$) of $\\langle \\dot M \\rangle<10^{-10} M_\\odot \\ {\\rm yr}^{-1}$ CVs in quiescence, allowing for discovery of the predicted large population of CVs with very low mass companions ($<0.1M_\\odot$) that are near the period minimum (Howell et al. 1997). Observations already show that the WD fixes the quiescent colors of these CVs and our calculations are useful for surveys in the field that were discussed at this meeting (e.g. 2DF, SDSS, see Marsh et al. 2001 and Szkody et al. 2001 contributions here), as well as HST CV searches in globular clusters. ", + "conclusions": "We have evaluated the action of compressional heating of accreting WD interiors. Most of the compressional energy release takes place in the accreted envelope, and is thermally communicated to the core. The maximum envelope mass is set by the unstable nuclear burning that causes a classical nova runaway and most likely expels the accreted mass. We have constructed equilibrium accretors which have constant core temperatures such that the heat lost from the core when the envelope is thin (i.e. right after the classical nova) is balanced by that regained when the envelope is thick. This equilibrium determines the $T_{\\rm eff}$ of the WD throughout the classical nova cycle. Our models agree with the observations of Dwarf Novae in deep quiescence and imply $\\langle \\dot M \\rangle\\approx10^{-10}M_\\odot$ yr$^{-1}$ just below the period gap and $\\langle \\dot M \\rangle\\approx10^{-9}M_\\odot$ yr$^{-1}$ just above the period gap for WD masses in the range $0.6$--$1.0M_\\odot$. Though our initial efforts have met with apparent success, there is still much to be done. First we need to investigate the relevant parameter ranges (e.g. $M$ and $\\langle \\dot M \\rangle$) within our initial scheme. The most critical parameter to vary is the metallicity of the accreted material, lowering to values relevant for globular cluster science. We must also survey our initial assumptions: \\begin{itemize} \\item WD excavation or accretion. The assumption that there was no net mass loss or gained by the WD through the classical nova cycle must be relaxed. This will allow for cooling of the WD due to adiabatic expansion after mass loss, or heating if it is increasing in mass. The large C/O fractions seen in some novae ejecta (Gehrz et al. 1998) might indicate that the WD is decreasing in mass at $\\langle \\dot M \\rangle$. \\item A self consistent accounting for the ignition masses, including varying the metallicity. Our initial work used Fujimoto's (1982) results for simplicity, but these are limited at low $\\langle \\dot M \\rangle$'s. A careful comparison to more modern nova calculations (e.g. Prialnik \\& Kovetz 1995) must be carried out. \\item Thermal evolution of the WD. We also need to concern ourselves with the secular change of $\\langle \\dot M \\rangle$ due to the decreasing companion mass or changing angular momentum losses (such as a drop in magnetic braking at the period gap). When this occurs on a timescale comparable to the WD thermal time, it is possible that the WD will not reach the steady-state solution we have assumed. This memory of previous higher $\\langle \\dot M \\rangle$ epochs could well allow the WD to be hotter than the equilibrium accretor. \\end{itemize} While these steps are unlikely to change our understanding of the compressional heating mechanism, each is essential for applying our work to the observations. As assumptions are relaxed and investigated, much more will be learned about both the state of CV systems and their evolution." + }, + "0112/astro-ph0112117_arXiv.txt": { + "abstract": "Magneto-convection simulations with two scenarios have been performed: In one, horizontal magnetic field is advected into the computational domain by fluid entering at the bottom. In the other, an initially uniform vertical magnetic field is imposed on a snapshot of non-magnetic convection and allowed to evolve. In both cases, the field is swept into the intergranular lanes and the boundaries of the underlying mesogranules. The largest field concentrations at the surface reach pressure balance with the surrounding gas. They suppress both horizontal and vertical flows, which reduces the heat transport. They cool, become evacuated and their optical depth unity surface is depressed by several hundred kilometers. Micropores form, typically where a small granule disappears and surrounding flux tubes squeeze into its previous location. ", + "introduction": "First, convection is driven from a very thin surface thermal boundary layer where radiation cools the fluid and produces the low entropy fluid which is pulled down by gravity and forms the cores of the cool downdrafts where most of the buoyancy work that drives the convection occurs. Second, convective flow is controlled by mass conservation. Fluid moving toward the surface in a stratified atmosphere, where density decreases upward, has to turn over and head back down in a distance the order of a scale height. Therefore upflows are diverging and fairly laminar, while downflows are converging and turbulent. As a result, only a small fraction of fluid at depth reaches the surface and there is little recycling of downflows into upflows. The horizontal scale of the upflows decreases toward the surface as the scale height decreases in proportion to the temperature. \\begin{figure}[htb] \\centerline{\\psfig{figure=b_h_tseq.eps,width=12cm}} \\caption[]{\\label{loops-t} {\\small Magnetic field loops rising through the solar surface and opening up. Time sequence is across then down. The upper boundary condition is a potential field. The domain is 6 $\\times$ 6 Mm $\\times$ 3 Mm deep.} } \\end{figure} Third, near the surface of the Sun hydrogen is partially ionized and most of the energy is transported as ionization energy which gets dumped at the surface when hydrogen becomes neutral. The downward kinetic energy flux is of order 10-15\\% of the enthalpy flux. The radiative flux is negligible below the surface. (In simulations with a large radiative flux in the interior. the temperature fluctuations are reduced and therefore larger velocities are required to carry a given flux. In such calculations there is a large downward kinetic energy flux.) Detailed results and comparisons with observations are given in Stein and Nordlund (1998, 2000). \\begin{figure}[htb] \\hbox{ \\epsfxsize=6.5cm \\epsffile{image_uz.contour_b.400G.547.2.eps} \\hspace*{0.3cm} \\vspace*{-0.2cm} \\epsfxsize=6.5cm \\epsffile{image_uzbot.contour_b0.400G.eps} \\vspace*{0.3cm} } \\hbox{ \\parbox[t]{6.6cm}{ \\caption[]{\\label{Vz-B-surf}{\\small Image of vertical velocity with magnetic field contours. Magnetic field is confined to the intergranular downflow lanes, but does not entirely fill them. The domain is 12 $\\times$ 12 Mm $\\times$ 3 Mm deep.} }} \\hspace{0.2cm} \\parbox[t]{6.6cm}{ \\caption[]{\\label{Vz-B-bot}{\\small Image of vertical velocity at depth 2.5 Mm with contours the surface magnetic field. The magnetic field collects in the mesogranular boundary downflows.} }} } \\end{figure} \\begin{figure}[htb] \\centerline{\\hbox{ \\epsfxsize=5.0cm \\epsffile{vizimage_bm.t=4.50.eps} \\hspace*{0.3cm} \\vspace*{-0.2cm} \\epsfxsize=6.5cm \\epsffile{b_hist-v-h.ps} \\vspace*{0.3cm} }} \\centerline{\\hbox{ \\parbox[t]{5.0cm}{ \\caption[]{\\label{B3D}{\\small 3D image of magnetic field. Magnetic flux tubes have a filamentary structure.} }} \\hspace{0.2cm} \\parbox[t]{7.0cm}{ \\caption[]{\\label{b_hist-v-h}{\\small The surface magnetic field has an exponential distribution for the case of a horizontal field advected in from the bottom, and it has a long tail with a bump near the value for pressure equilibrium with its surroundings for the case of a uniform vertical field.} }} }} \\end{figure} ", + "conclusions": "" + }, + "0112/astro-ph0112321_arXiv.txt": { + "abstract": "The KPNO International Spectroscopic Survey (KISS) is an objective-prism survey for extragalactic emission-line objects. It combines many of the features of previous slitless spectroscopic surveys with the advantages of modern CCD detectors, and is the first purely digital objective-prism survey for emission-line galaxies. Here we present the first list of emission-line galaxy candidates selected from our blue spectral data, which cover the wavelength range 4800 -- 5500 \\AA. In most cases, the detected emission line is [\\ion{O}{3}]$\\lambda$5007. The current survey list covers a one-degree-wide strip located at $\\delta$ = 29$\\arcdeg$~30$\\arcmin$ (B1950.0) and spanning the right ascension range 8$^h$~30$^m$ to 17$^h$~0$^m$. An area of 116.6 deg$^2$ is covered. A total of 223 candidate emission-line objects have been selected for inclusion in the survey list (1.91 deg$^{-2}$). We tabulate accurate coordinates and photometry for each source, as well as estimates of the redshift, emission-line flux, and equivalent width based on measurements of the digital objective-prism spectra. The median apparent magnitude of the sample is B = 18.2, and galaxies with redshifts approaching z = 0.1 are detected. The properties of the KISS emission-line galaxies are examined using the available observational data, and compared to previous surveys carried out with photographic plates as well as with the H$\\alpha$-selected portion of KISS. ", + "introduction": "The KPNO International Spectroscopic Survey (KISS) is an ongoing objective-prism survey which targets the detection of large numbers of extragalactic emission-line sources. KISS attempts to build upon previous, extremely fruitful surveys for these types of objects, which have been responsible for the cataloging of a large proportion of the known starburst galaxies, Seyfert galaxies, and QSOs. Our survey method is similar to many of these previous surveys, which have been carried out with Schmidt telescopes and photographic plates (e.g., Markarian 1967, Smith \\etal 1976, MacAlpine \\etal 1977, Pesch \\& Sanduleak 1983, Wasilewski 1983, Markarian \\etal 1983, Zamorano \\etal 1994, Popescu \\etal 1996, Surace \\& Comte 1998, Lipovetsky \\etal 1998). The fundamental difference between KISS and these previous surveys is that we utilize a CCD as our detector. There are several obvious advantages of CCDs over photographic plates for this type of survey, including much higher quantum efficiency, lower noise, good spectral response over the entire optical portion of the spectrum, and large dynamic range. In addition, CCDs enable us to use automated selection methods to detect emission-line galaxies (ELGs), and allow us to quantify the selection function and completeness limit directly from the survey data. With the advent of large format CCDs in the past decade, the large areal coverage provided by the wide-field imaging capability of Schmidt telescopes makes digital surveys like KISS possible. The combination of increased depth and large areal coverage leads to substantial improvements compared to the previous photographic surveys listed above. The primary goal of KISS is to produce a high-quality survey whose selection function and completeness limits can be accurately quantified so that the resulting catalog of ELGs will be useful for a broad range of studies requiring statistically complete galaxy samples. We also want to reach substantially deeper than previous objective-prism surveys. A complete description of the survey method employed for KISS is given in the first paper in this series (Salzer \\etal 2000, hereafter Paper I). KISS is a {\\it line-selected} survey, meaning that the objective-prism spectra are searched for the presence of an emission feature. As described in Paper I, the first survey strip was observed in two distinct spectral regions. One covered the blue portion of the optical spectrum (4800 -- 5500 \\AA), while the second covered the wavelength range 6400 -- 7200 \\AA\\ in the red part of the spectrum. The first red survey list is given in Salzer \\etal (2001, hereafter KR1). The current paper presents the initial KISS list of [\\ion{O}{3}]-selected ELG candidates. The format of the current paper follows closely that of KR1. In addition to listing the ELGs, we provide substantial observational data for each object. This includes accurate photometry and astrometry for each source, as well as estimates of each galaxy's redshift, line flux, and equivalent width. These data are used to examine the properties of the KISS ELGs in Section 4. ", + "conclusions": "We present our second list of emission-line galaxies discovered as part of the KPNO International Spectroscopic Survey (KISS). This is the first list of [\\ion{O}{3}]-selected ELGs. A total of 223 galaxies are included in this survey list, which covers an area of 116.6 deg$^2$. With a surface density of 1.91 ELGs per deg$^2$, the blue portion of KISS finds more than 19 times the number of AGNs and starbursting galaxies per unit area than the Markarian survey, and nearly 4 times the number found by the UM survey, which used a similar selection method. An additional 91 ELG candidates detected with a slightly lower level of statistical significance are cataloged in a supplementary list (see appendix). The main advantages of KISS over previous photographic surveys are the combination of increased depth and wavelength coverage afforded by the use of a CCD as the survey detector. The digital nature of the survey data, which includes both imaging and spectral observations, means that a great deal of information is available for each object in the sample. In addition to tabulating the ELG candidates, Table 2 also includes accurate astrometry, B and V photometry, estimates of the redshift, and emission-line strength information for all objects in the catalog. This allows us to investigate the properties of the survey constituents without the need for detailed follow-up observations. We illustrate the distributions of apparent magnitudes, colors, emission-line strengths, redshifts, and absolute magnitudes for the full sample of KISSB ELGs. The new catalog of ELGs is found to have an apparent magnitude distribution very similar to that of KISSR, with a median B magnitude of 18.17. Not surprisingly, the typical colors of the KISSB galaxies are somewhat bluer than those for KISSR. Based on the distribution of measured emission-line equivalent widths, we estimate that KISSB detects most ELGs in the survey area with [\\ion{O}{3}] EW $>$ 30 \\AA. The luminosity distribution is skewed toward lower luminosities (median M$_B$ = $-$18.04) than KISSR. This last item is an interesting result, and is due to the use of [\\ion{O}{3}]$\\lambda$5007 rather than H$\\alpha$ as the primary emission line used to select the survey constituents. While KISSR and KISSB are in all other aspects similar surveys, selecting the ELGs via the two different emission lines changes the make-up of the respective samples in a major way. We also compare the properties of the KISS galaxies with those of previous surveys (e.g., Markarian, Case, UM, UCM) in order to evaluate the relative strengths of each, as well as to better understand how their selection functions help to shape the nature of the resulting samples. The KISSB galaxies are fainter (by 1.3 to 2.6 magnitudes on average) than the galaxies detected in this representative sample of photographic surveys. The physical characteristics of the KISSB ELGS (e.g., color and luminosity) compare most closely with those of the UM survey, which is also [\\ion{O}{3}] selected. Like the UM survey ELGs, the KISSB galaxies possess a wide range to physical properties and activity levels, and include both luminous starbursting and Seyfert galaxies, as well as many intermediate- and low-luminosity star-forming systems. We note that, despite the wealth of information available for the KISS ELGs from the survey data alone, follow-up spectra are still required in order to confirm each source as a {\\it bona fide} emission-line galaxy. In addition, these spectra are necessary for providing more accurate redshifts as well as emission-line strengths in order to classify each ELG by their activity type (i.e., AGN {\\it vs.} starburst). To date, we have obtained follow-up spectra for slightly more than half of the KISSB candidates; 119 of 123 KISSB candidates with spectra are found to be real ELGS (97\\%). Additional follow-up spectroscopic observations are in progress for large subsets of the KISS ELG catalogs. Additional lists of ELG candidates are currently being prepared for publication (e.g., Gronwall \\etal 2002b), and observational data continue to be obtained for new survey areas. All future observations will focus on the H$\\alpha$ portion of the spectrum, for the reasons described in the text. The overall goals of the KISS survey are to cover roughly 300 sq. deg. of sky, and to catalog in excess of 5000 ELG candidates." + }, + "0112/hep-ph0112278_arXiv.txt": { + "abstract": " ", + "introduction": "One of the very attractive arguments in favour of supersymmetry(SUSY) is that it provides a natural solution to the dark matter problem. In R-parity conserving supersymmetric models there exists a neutral stable particle, the lightest supersymmetric particle(LSP), which could constitute the cold dark matter in the universe. As there are strong constraints on stable charged particles, only if the LSP is a neutralino or a sneutrino could SUSY provide a suitable dark matter candidate. Cosmological constraints however as well as direct searches disfavour the sneutrino \\cite{sneutrino} leaving the neutralino as the preferred candidate for dark matter. The contribution of neutralinos to the relic density is however very model dependent and varies by several orders of magnitude over the whole allowed parameter space of the MSSM. The relic density then imposes stringent constraints on the parameters of the MSSM often favouring solutions with light supersymmetric particles. It has been known for some time that coannihilation processes where the LSP interacts with only slightly heavier sparticles can significantly reduce the estimate of the relic density, leading to acceptable values even with a rather heavy sparticle spectrum. In principle these coannihilations can occur with any supersymmetric particle\\cite{Griest}. The importance of the coannihilation channels were emphasized before both for gauginos \\cite{Yamaguchi,EdsjoGondolo}, sleptons\\cite{Ellis-coann,GLP} or stops \\cite{abdelstop,Ellisstop}. Here, {\\em ALL} channels will be included. In some regions of the parameter space the neutralinos annihilate so fast that they cannot constitute the only source of dark matter. This happens for example when the masses are such that the neutralinos can annihilate through a s-channel Higgs resonance \\cite{Ellis-Higgs,DreesNojiri} or a s-channel Z resonance\\cite{g-2nous}. In this context the inclusion of the relic density constraints, by giving a handle on the supersymmetric particle spectrum, has important consequences both for studies of SUSY at colliders and in astroparticle experiments. There exist many calculations of the relic density in supersymmetry, using various approximations both in the evaluation of cross-sections and in solving the density equation \\cite{Ellis-Higgs},\\cite{neutdriver},\\cite{Darksusy},\\cite{Leszek},\\cite{Torino}, \\cite{BaerBrhlik}. Among these, {\\tt Neutdriver}\\cite{neutdriver}and {\\tt DarkSusy}\\cite{Darksusy} are publicly available. Our purpose is to provide a tool that evaluates with high accuracy the annihilation cross-sections even in regions near poles and thresholds, that is both flexible and upgradable and that goes beyond {\\tt DarkSusy} as far as the calculation of the relic density is concerned. The main characteristics of this program, called \\micro, are \\begin {itemize} \\item{} Complete tree-level matrix elements for all subprocesses \\item{} Includes all coannihilation channels with neutralinos, charginos, sleptons, squarks and gluinos. \\item{} Loop-corrected Higgs masses and widths \\item{} Speed of calculation \\end{itemize} All calculations of cross-sections are based on \\comphep\\cite{comphep}, an automatic program for the evaluation of tree-level Feynman diagrams. We follow the method proposed by \\cite{GondoloGelmini} for the calculation of the relic density with its generalization to the case of coannihilations \\cite{EdsjoGondolo}. We still rely on approximations for the solution of the relic density equations and the determination of the freeze-out temperature, since this allows to significantly increase the speed of the program and proves to be very useful when scanning over a large parameter space. This, together with the fact that we have included sfermion coannihilation channels as well as one-loop corrections to the Higgs width constitute the main differences with {\\tt DarkSUSY}. Although we will generally assume that the neutralino is the LSP, \\micro~ can be used to compute the relic density with any supersymmetric particle as the LSP, in particular the sneutrino. This is because all (co-)annihilation of any pairs of supersymmetric particles into any pairs of standard model or Higgs particles are included. The program for the relic density calculation described here is contained in a package that lets the user choose between weak scale parameters or parameters of SUGRA models as input parameters. The latter is achieved through a link with \\isasugra\\cite{ISASUGRA}. The calculation of the Higgs masses are done with {\\tt FeynHiggsFast}\\cite{FeynHiggs}. Loop QCD corrections to the Higgs partial widths into fermions are extracted from {\\tt HDECAY}\\cite{HDECAY}. In addition we provide subroutines that calculate various constraints on the MSSM parameters: direct limits from colliders, $\\Delta\\rho$, $b\\to s\\gamma$ and $(g-2)_\\mu$. All these constraints can be updated or replaced easily. The total number of processes which can contribute to the relic density exceeds 2800. However, due to a strong Boltzmann suppression factor, only processes with SUSY particles close in mass to the LSP are relevant for the calculation. Therefore in most cases only a small fraction of the available processes are needed. In principle, compilation of the full set of subprocesses is possible, but such a program would be huge and could not be distributed easily. To avoid this problem, we include in our package the program \\comphep\\cite{comphep} which generates, while running, the subprocesses needed for a given set of MSSM parameters. The generated code is linked during the run to the main program and executed. The corresponding ``shared\" library is stored on the user disk space and is acessible for all subsequent calls, thus each process is generated and compiled only once. Such approach can be realized only on Unix platforms which support dynamic linking. The paper is organized as follows. After we summarize the important equations for the calculation of the relic density, we give a short description of the parameters of the supersymmetric model. A description of the package follows. Finally we present some results and comparisons with another program in the public domain, \\darksusy. ", + "conclusions": "The package \\micro~ allows to calculate the relic density of the LSP in the most general MSSM with $R_P$ conservation. The package is self-contained safe for the {\\tt ISASUGRA} package that is required when using the SUGRA option. All possible channels for coannihilations are included and all matrix elements are calculated exactly at tree level with the help of {\\tt CompHEP}. Loop corrections for the masses of Higgs particles (two-loop) and the width of the Higgs (QCD one-loop) are implemented. Good agreement with existing calculations is found when identical set of channels are included. Future versions will include loop corrections to neutralino masses. Even though these corrections are only a few GeV's they can alter significantly the calculation of the relic density when there is a near mass degeneracy with the next to lightest supersymmetric particle that contributes to a coannihilation channel \\cite{benchmark}. Although the loop processes are in general small, we will include $\\chi\\chi\\ra\\gamma\\gamma,\\gamma Z, gg$ in an update of \\micro." + }, + "0112/astro-ph0112429_arXiv.txt": { + "abstract": "{% We study the evolution of embedded protoplanets in a protostellar disk using very high resolution nested-grid computations. This method allows us to perform global simulations of planets orbiting in disks and, at the same time, to resolve in detail the dynamics of the flow inside the Roche lobe of the planet. The primary interest of this work lies in the analysis of the gravitational torque balance acting on the planet. For this purpose we study planets of different masses, ranging from one Earth-mass up to one Jupiter-mass, assuming typical parameters of the protostellar disk. The high resolution of the method allows a precise determination of the mass flow onto the planet and the resulting torques. The obtained migration time scales are in the range from few times $10^4$ years, for intermediate mass planets, to $10^6$ years, for very low and high mass planets. Typical growth time scales depend strongly on the planetary mass, ranging from a few hundred years, in the case of Earth-type planets, to several ten thousand years, in the case of Jupiter-type planets. ", + "introduction": "During the past five years radial velocity studies have allowed the detection of planetary companions around other main-sequence stars. Until now about sixty so-called ``extrasolar planets'' have been discovered, which orbit their stars within a distance of a few \\AU. A recent catalog of extrasolar planets, including their orbital characteristics, is provided by Butler et al. (\\cite{butler2001}) and up-to-date versions can be found at \\texttt{http://www.obspm.fr/encycl/encycl.html} and \\texttt{http://exoplanets.org/}, maintained by Jean Schneider and the Department of Astronomy at UC Berkeley, respectively. In contrast to the solar system, these new planets display quite different orbital properties that challenge the accepted formation scenario for solar planets. The major differences are their high minimum masses (up to $17$ Jupiter-masses), their proximity to the central star (a fraction of the Sun-Mercury distance) and their high eccentricities (up to $0.7$). One of the main problems to deal with is the very close distance of massive planets to their parent star. The formation of Jupiter-type planets at these locations is, on theoretical grounds, very unlikely. First of all, from purely geometrical arguments, the matter reservoir of the surrounding disk is too little so that a planet could hardly accrete its mass. Second, the temperatures within the disk are too high for a rocky core to condense easily. For these reasons it is generally believed that planets have formed from disk material further out, at distances of several \\AU\\ from the star, and have then migrated to their present positions. This radial motion of the planet through the disk is primarily caused by gravitational torques acting on the planet. The presence of the planet in the disk disturbs the disk gravitationally, creating spiral density wave perturbations, which emanate from the planet through the disk. Hence, the disk is no longer axisymmetric which results in a net torque on the planet. The sign and magnitude of the vertical component of the torque determines the direction and efficiency of the radial migration. While initial fully non-linear hydrodynamical numerical computations of embedded planets assumed a fixed circular orbit of the planet (Kley \\cite{kley1999}; Bryden et al. \\cite{bryden1999}; Lubow et al. \\cite{lubow1999}), more recent simulations took into account the back reaction of the disk and allowed for a change in the parameters of the planetary orbit (Kley \\cite{kley2000}; Nelson et al. \\cite{rnelson2000}). For a Jupiter-mass planet and typical parameter values for the disk, the obtained orbital decay time is about $10^5$ years, which agrees reasonably well with previous estimates based on analytic linear theories (Goldreich \\& Tremaine \\cite{gt1980}; Ward \\cite{ward1997}). The majority of the computations, performed so far, have used a single grid which resolves the Roche lobe of a Jupiter-mass planet only with very few grid cells. Recently, Cieciel\\c{a}g et al. (\\cite{ciecielag2000a}, \\cite{ciecielag2000b}) used an Adaptive Mesh Refinement method to resolve the immediate surroundings of the planet, but they didn't give any estimate of the mass accretion rate and magnitude of the gravitational torque. On the other hand, Armitage (\\cite{armitage2001}) reduced the overall simulated region achieving a better resolution. However, also in this case the Roche lobe is only resolved by a few grid cells because of the low mass of the investigated planet. In this paper we aim at the structure and dynamics of the gas flow in the close vicinity of the planet, while performing global disk simulations. In order to obtain the necessary high spatial and temporal resolution, we use a nested-grid formalism which allows an accurate computation of the mass flow onto the planet and the acting torques. In the next section we layout the physical model followed by a description of the numerical method (\\Sect{Sect:NM}). We describe the setup of the various numerical models in \\Sect{Sect:GMD}. The main results are presented in \\Sect{Sect:results} and our conclusions are given in \\Sect{Sect:conclusions}. ", + "conclusions": "\\label{Sect:conclusions} A number of numerical simulations concerning disk-planet interactions have been performed to get new insights into the scenario of the joint evolution of protoplanets and their environment. They have confirmed analytical theories for gap formation and planet migration. However, many open questions still remain. The most important unsolved issue is the influence of the ambient gas on the dynamical evolution of a planet. Another one is the way disk-planet interaction changes when small planets, in the mass range of Neptune and Earth, are considered. We began to investigate in both directions by means of a nested-grid technique, which is particularly suitable for treating these problems. The main asset of this numerical scheme is the possibility of achieving, locally, a very high spatial and temporal resolution. With such a method we are able to resolve very accurately both, the inner parts of the Hill sphere of the planet \\emph{and} the global structure of the disk, because we treat the whole azimuthal extent of the disk. Concerning the issues mentioned above, with the present paper we tackled some outstanding problems concerning the growth and migration of protoplanets, covering a range from one Jupiter-mass down to one Earth-mass. Thus, even though we do not include the detailed energetic balance of the planetary structure, which is still beyond present day computer facilities, this study represents a definite improvement in the determination of the torque balance on protoplanets. Our main achievements can be summarized as follows. \\begin{enumerate} \\item Inside the Hill sphere of the planet a circumplanetary disk forms. Within it strong spiral density waves develop, if the planet mass is larger than $\\sim 5$ \\MEarth. These waves assume the shape of a two-arm pattern. The two spiral arms are slightly asymmetric with respect to the planet. For decreasing planet masses, they stretch and shorten. Matter is observed to pile up at the location of the planet, generating a very high density zone (we named as density ``core''), which might represent its primordial gaseous envelope. \\item Nearby material exerts positive torques on the planet, slowing down, considerably in some cases, its inward migration. Most of these torques arise from corotation regions, i.e. from gas lying on the planet's orbit. Analytical models about migration do not account for them. This can be one of the reasons why our estimates of the migration time scales give somewhat higher values than those predicted by such theories. \\item Within a distance of $\\sim 0.1$ \\Rhill\\ from the planet, the point-mass approximation becomes too restrictive and maybe not appropriate. Therefore, the structure of the planet should be also taken into account over such a length scale. This is absolutely necessary if one wants to evaluate how much of the angular momentum, transferred by closely orbiting matter, is conveyed to the spin of the planet rather than to its orbital angular momentum. \\item The Keplerian rotational regime of circumplanetary disks is affected by spiral perturbations. Just as for the mass density, the more massive the planet is, the stronger such perturbations are. Gas material, passing through the spiral fronts, is deflected towards the planet. Instead, in the inter-arm regions it moves away from it. This is in analogy with the spiral wave theory in galaxies. Around Earth-mass planets, the rotation of the gas is very slow if compared to the Keplerian rotation. In fact, in this particular case, the density core has nearly a hydrostatic structure. \\item The mass accretion rate, as a function of the mass of the planet, has a maximum around $\\Mp = 0.5$ \\MJup. As long as $\\Mp \\lesssim 0.2$ \\MJup, the growth time scale of a planet increases, approximatively, as $\\Mp^{2/3}$. For more massive planets, it increases roughly as $\\Mp^{4/3}$. Such a dependence may contribute to limit the size of a massive planet. \\end{enumerate} Since we just started to explore these new grounds, each of the items above may deserve a more specific and dedicated study. Next efforts should be devoted to refine the physical model, especially in the vicinity of the planet. Henceforth, some of the future developments could be: \\begin{itemize} \\item including an energy equation, by implementing an approximate treatment of radiative transfer and viscous dissipation; \\item improving the equation of state by using an alternative form which accounts for the planet's structure; \\item evaluating possible effects due to the two-dimensional approximation of the disk, via three-dimensional simulations. \\end{itemize}" + }, + "0112/astro-ph0112103_arXiv.txt": { + "abstract": "We obtained new near-infrared images of the prototypical pre-main sequence triple system T\\,Tau, as well as the first resolved medium-resolution spectra of the close pair {\\tts}. At the time of our observations, the tight binary had a 13\\,AU projected separation and showed significant motion since its discovery, three years before. The orbit cannot be strongly constrained yet, but the observed motion of {\\ttb} with respect to {\\tta} suggests that the system is at least as massive as {\\ttn} itself. This may indicate that {\\ttn} is not the most massive star in the system. The spectrum of {\\tta}, which is totally featureless except for a strong {\\brg} emission line, identifies this component with the ``infrared companion'', whose exact nature remains obscure but may be the consequence of it being the most massive component of the system. Contrasting sharply with {\\tta}, the spectrum of {\\ttb} shows numerous photospheric features consistent with an early-M spectral type. The presence of a strong {\\brg} emission line and of a significant veiling continuum classifies this object as a deeply embedded T\\,Tauri star. From these observations, we conclude that both components of {\\tts} are embedded in their own dense circumstellar cocoon of material, which are probably fed by a much more extended structure. ", + "introduction": "T\\,Tau has long been identified as one of the brightest low-mass pre-main sequence objects and, as such, has been considered as the prototype for the T\\,Tauri class of objects \\citep{joy45}. Its notable properties include a significant infrared continuum excess thought to be related to the presence of an accretion disk around this million year-old object \\citep{bertout88}, a bright nebulosity caused by photon scattering on the inner wall of a cavity emptied by a strong polar outflow \\citep{stapelfeldt98}, a substantial circumstellar disk resolved through interferometric imaging at radio wavelengths \\citep{akeson98} and a limited interstellar extinction \\citep[$A_V\\sim1.5\\,$mag,][]{ghez91}. Early near-infrared high-angular resolution observations revealed a close \\citep[0\\farcs7, or 100\\,AU at $d=140\\,$pc,][]{elias78} companion in this system \\citep{dyck82}, which raised the possibility that T\\,Tau is not so prototypical of its class. From subsequent various high-angular resolution surveys \\citep{ghez93,leinert93,simon95} however, it rapidly became clear that binarity is not a rare property, re-establishing T\\,Tau as a prototype of pre-main sequence solar-type stars once again. The companion to T\\,Tau, known as \\tts, is a very peculiar one. While it is undetected in visible images obtained first by speckle interferometry \\citep{gorham92} and more recently by HST/WFPC2 \\citep{stapelfeldt98}, implying a huge flux ratio at those wavelengths ($\\Delta V>9.5$), its bolometric luminosity is about twice as large as that of {\\ttn}. Indeed, {\\tts} dominates the system flux at wavelengths $\\gtrsim2\\,\\mu$m and its spectral energy distribution peaks around 3\\,$\\mu$m, a significantly longer wavelength than normal for T\\,Tauri stars \\citep{ghez91}. Near- to mid-infrared monitoring has further revealed that {\\tts} is strongly variable, up to 2~mag over timescales of a few months \\citep{ghez91}. Finally, its near-infrared spectrum is featureless around 2\\,$\\mu$m at a resolution of $R\\sim760$, with the exception of a large {\\brg} emission line and a much weaker H$_2$ line \\citep{beck01}. Together with a handful of other companions to T\\,Tauri stars, {\\tts} has been classified as an ``infrared companion'' (IRC). As reviewed by \\cite{koresko97}, the peculiar properties of these companions has generated a wide range of theories regarding their nature, which include an embedded intermediate- to high-mass protostar, a planetary object embedded in the disk surrounding their optical companion or a strongly accreting FU\\,Ori-like object among others possibilities. The very nature of {\\tts} and other IRCs is still debated as none of these explanations fit all the observational data. Increasing the complexity of the T\\,Tau system, \\cite{koresko00} identified a very close companion to {\\tts} through speckle holography. In the discovery images, the companion was projected a mere 7\\,AU (0\\farcs05) away from the IRC. Such a small separation raised new questions about the {\\tts} system: are both components classifiable as IRCs? If embedded, are they inside the same cocoon of material? Can the additional companion be used to clarify the nature of IRCs? While previous high-angular resolution observations of this system never resolved the IRC \\citep[e.g.,][]{simon96}, follow-up observations by \\cite{kohler00} confirmed the presence of a close companion. However, they noted significant changes in the binary separation and position angle, which cannot be accounted for by measurement uncertainties, but rather suggests that the two objects are bound with an orbital period on the order of 10--20\\,yrs. In this paper, we present adaptive optics images obtained with the Keck telescope as well as the first spatially resolved $K$-band spectra of the {\\tts} system. Our observations and data reduction procedures are described in \\S\\,\\ref{sec:obs-red} and our main results are presented in \\S\\,\\ref{sec:results}. Implications on the nature of the IRC and its companion are presented in \\S\\,\\ref{sec:discus}. Finally, \\S\\,\\ref{sec:conclu} summarizes our findings. ", + "conclusions": "\\label{sec:conclu} We obtained new $H$ and $K$ images of the T\\,Tau triple system with the adaptive optics system at the 10\\,m Keck telescope, as well as the first resolved medium-resolution ($R\\sim3500$) $K$-band spectra of the two components of {\\tts}, identified so far as an IRC to {\\ttn}. Our data reveal that {\\ttb} is a heavily extincted, actively accreting M1 pre-main sequence object, i.e. a rather normal T\\,Tauri star. On the other hand, {\\tta}, which dominates the system in the near-infrared, is indeed identified as the IRC: its spectrum shows no feature except for a significant {\\brg} emission line. Unlike previous observations, we do not detect the 2.12\\,$\\mu$m H$_2$ emission line in either of the two components. The tight {\\tts} binary shows significant motion since its first detection, almost three years prior to our observations. Although it is not yet inconsistent with a constant linear motion, it seems quite unlikely that this is because either one of the two objects is a background object, since they both show strong signs of T\\,Tauri-like activity. Interpreting the observed motion as that of two stars orbiting each other, and assuming that the orbit is inclined and circular, we derive rough estimates for the orbital parameters. The total system mass and orbital periods are on the order of $5\\,M_\\odot$ and $25$\\,yrs, although we emphasize that the uncertainty on the actual orbit shape is not taken into account here. Although {\\tts} might well be located behind the circumstellar disk of {\\ttn}, we argue here that this is {\\it not} the main reason for the peculiarities of {\\tts}, because of the strikingly different properties of {\\tta}, which is now identified as the IRC, and {\\ttb} as well as from an analysis of the optical depth of the circumprimary disk. Both components of {\\tts} are likely embedded in their own circumstellar material, which could be either an almost edge-on disk or a dense spheroidal envelope. Because of tidal truncation due to the tight binary's orbit, this material must be confined to within only a few AU of the stars, making it too small to be resolved by current observations. Furthermore, this effect also explains the non-detection of thermal flux in the radio domain, as most of the material located in the corresponding range of distance from the stars has been swept away. In any event, the subsistence of dense envelopes on a few AU scale around both components of {\\tts} for several $10^5$\\,yrs requires that they are being replenished by material coming from a larger scale reservoir such as a vast envelope surrounding the whole system. The presence of such a structure was already suspected from the analysis of the spectral energy distribution from the system. Following the orbit of {\\tts} for a few more years will bring a definitive conclusion regarding the possibility that the system is simply a projected, unphysical one and provide a much better estimate of the orbit shape and parameters if the two stars are bound. Furthermore, high-spectral resolution data will yield radial velocities, which are very powerful tools to better constrain the orbit, and potentially reveal for the first time some photospheric features, which would greatly help understanding the nature of this peculiar object. We also note that both components of {\\tts} are prime targets for the upcoming long-baseline interferometry experiments, as the environment of these stars can only be resolved through these techniques." + }, + "0112/astro-ph0112273_arXiv.txt": { + "abstract": "{ We investigate the behaviour of cosmological baryons at low redshifts $z \\la 5$ after reionization through analytic means. In particular, we study the density-temperature phase-diagram which describes the history of the gas. We show how the location of the matter in this $(\\rho,T)$ diagram expresses the various constraints implied by usual hierarchical scenarios. This yields robust model-independent results which agree with numerical simulations. The IGM is seen to be formed via two phases: a ``cool'' photo-ionized component and a ``warm'' component governed by shock-heating. We also briefly describe how the remainder of the matter is distributed over galaxies, groups and clusters. We recover the fraction of matter and the spatial clustering computed by numerical simulations. We also check that the soft X-ray background due to the ``warm'' IGM component is consistent with observations. We find in the present universe a baryon fraction of 7\\% in hot gas, 24\\% in the warm IGM, 38\\% in the cool IGM, 9\\% within star-like objects and, as a still un-observed component, 22\\% of dark baryons associated with collapsed structures, with a relative uncertainty no larger than 30\\% on these numbers. ", + "introduction": "\\label{Introduction} As is well known, the mass of baryons embedded within stars or galactic disks in the current universe is quite small since it yields a baryonic parameter $\\Omega_{\\rm gal} \\sim 0.002 - 0.006$ (e.g., \\cite{Fuk1}) while standard nucleosynthesis calculations give $\\Ob \\simeq 0.045$ (e.g., \\cite{Tyt1}). Therefore, most of the baryonic matter should lie in the intergalactic medium. This agrees rather well with the fact that at higher redshift one observes a large amount of mass in the Lyman-$\\alpha$ forest which consists of moderate density fluctuations ionized by the background UV flux emitted by distant galaxies. Thus, one gets $\\Omega_{\\rm Ly \\alpha} \\sim 0.01-0.05$ at $z\\sim 3$ (e.g., \\cite{Fuk1}). However, as noticed in \\cite{Cen1} the mass within the Lyman-$\\alpha$ forest decreases with time and at $z=0$ summing over all observed contributions one obtains $\\Ob \\la 0.011$ which falls short of the required total baryonic mass. Hence at $z=0$ a large part of the baryons must lie in a new intergalactic component beyond the usual Lyman-$\\alpha$ forest clouds. As argued in \\cite{Cen1} and \\cite{Dave2} this could be part of a ``warm'' phase of the intergalactic medium (IGM), with temperatures in the range $10^5 < T <10^7$ K. The latter conclusion was reached from numerical simulations. In this article, we reconsider this problem in order to derive the properties of the IGM by analytic means. In particular, we wish to investigate whether one can understand this behaviour in a quantitative manner from robust, model-independent, arguments. First, in Sect.\\ref{The phase-diagram of cosmological baryons} we study the $(\\rho,T)$ phase-diagram of cosmological baryons. While the Lyman-$\\alpha$ forest is described by a well-defined Equation of State the ``warm'' IGM component shows a broad scatter (e.g., \\cite{Dave1}, \\cite{Dave2}) since its temperature depends through shock-heating on the neighbouring gravitational potential which is a stochastic field. Nevertheless, we show that it is constrained to lie in a well-defined domain in the $(\\rho,T)$ plane, and determine its average location in this plane, which may be considered as the ``Equation of State'' of the ``warm'' IGM. We also give the location in this diagram of galaxies, groups and clusters. Next, in Sect.\\ref{Distribution of matter} we use our results to compute the redshift evolution of the fraction of matter enclosed within the different phases. Then, in Sect.\\ref{Warm IGM two-point correlation} we estimate the two-point correlation function and the clumping factor of the ``warm'' IGM. Finally, in Sect.\\ref{Soft X-ray background} we check that the X-ray background emitted by the ``warm'' component agrees with observations. ", + "conclusions": "In this article we have shown that the $(\\rho,T)$ phase-diagram of cosmological baryons can be understood in very simple terms. It is strongly model-independent since the relevant scales are set by atomic physics (hydrogen ionization) and the basic cosmological setting (the rate of expansion and the scale which marks the transition to the non-linear regime). Building a bridge between the modeling of galaxies, groups and clusters, as well as Lyman-$\\alpha$ absorbers, we have constructed in this paper a phase-diagram which includes also the IGM, and thus provides a {\\it phase-diagram of all baryons}. In this diagram, all these objects form well-defined populations. Firstly, we have distinguished a ``cool'' IGM phase corresponding to the Lyman-$\\alpha$ forest. It follows a well-defined equation of state on the $(\\rho,T)$ plane. We have shown that the properties of this gas do not depend much on its previous history because of the form of the recombination coefficient. This explains why the scatter of this component obtained in numerical simulations is quite small. The fraction of baryons in this components, as well as its evolution with redshift is seen to be consistent with earlier modelling of the frequency of occurrence of the Lyman-$\\alpha$ lines. We however obtain (consistently with the numerical simulations) 38\\% of the baryons in the cool IGM at z=0. This may be a little high with a mass fraction closer to ~30\\% being consistent with observations. This must be considered as being within the error bars of our analytical calculation (and the error bars of the simulations!). Secondly, we find that a ``warm'' IGM phase is formed by the gas which has been shock-heated to larger temperatures $T \\sim 10^4 - 10^7$ K as non-linear gravitational structures appear. The dependence on the stochastic gravitational potential entails a broad scatter in the $(\\rho,T)$ plane. However, we have explained that robust constraints only allow a closed region in the $(\\rho,T)$ phase-diagram for this component. We have also defined a simple curve which represents its mean behaviour and plays the role of the ``warm'' IGM Equation of State. The latter is found to be consistent with the outcome of numerical simulations (\\cite{Dave2}), but differs from the simple power-law fit of a mean curve done there. This is because its locus is based on physical considerations, which we have thoroughly justified. The ``warm'' IGM is seen to be due to large structures on their way of collapsing under the action of gravity. The quite simple picture behind this model is that the collapse of an object turning non-linear (that is with average overdensity above $5$) induces shocks which heat patches of much smaller size, to a temperature of the order of the locally available kinetic energy of the collapse, creating local conditions within hot spots close to hydrostatic equilibrium. We are able to estimate this size, as well as the temperature of the latter, and to show they are in agreement with the simulations. Definitely, these objects are not just Lyman-{$\\alpha$} absorbers at larger densities. The hot gas which lies within clusters and galactic halos, on the other hand, corresponds to the high-density ``continuation'' (e.g., $\\delta \\ga 200$) of the ``warm'' IGM curve, in a region of phase-space where virialization insures that all the available gas has been gravitationally heated. Next, we find two quite different regimes for the ``warm'' IGM. At high redshift, some of the ``warm'' component enters the cooling region of the $(\\rho,T)$ plane: this gives rise to galactic disks and stars. Some of the baryonic matter, also, should lie on the low temperature branch of this cooling region, in agreement with the results of numerical simulations (\\cite{Dave1}). At low redshift, the ``warm'' IGM no longer goes into the cooled region, and it provides the origin of the hot gas in clusters. {\\it The evolution with redshift of the baryon phase-diagram provides a natural confirmation that at high $z$ collapsed halos form galaxies while at low redshift they build groups or clusters}. The warm phase occupies 24\\% of the baryon fraction at $z=0$ in our analytic model. Then, we have checked that our results for the fraction of matter enclosed within the various phases and the two-point correlation function of the ``warm'' IGM component agree with numerical simulations (e.g., \\cite{Dave1} and \\cite{Dave2}). This confirms the validity of our analysis. Note however that the latter relies on simple physical considerations and it is independent of the findings of the former. Besides, the soft X-ray background due to the ``warm'' IGM is consistent with the upper bound set by observations. Our prediction actually is not far from the observational limit: this offers the prospect of measuring this X-ray emission in future observations. We predict that approximately 60\\% of the baryons are accounted for at z=0 in the cool/warm phases of the intergalactic medium. The remaining baryons (i.e. 40\\%) are in collapsed structures. With a total of $\\Omega_b \\approx 0.045$, this amonts to $\\Omega_{\\rm bcoll} \\approx 0.018$. The observed luminous component of the baryons (stars, remnants and gas) is estimated to be about $\\Omega_{\\rm gal} \\approx 0.004 $, as noted in the introduction, and the hot gas in clusters represents $\\Omega_{\\rm ghot} \\approx 0.003$. Therefore, there remains an ``unidentified'' or dark baryon component, $\\Omega_{\\rm bdm} \\approx 0.011$. This is at least twice the stellar-like component but only about 1/4 of the ``identified'' baryons (IGM, stellar components and very hot gas). In our model this dark baryonic matter corresponds to (cool ?) gas within galactic halos and groups, which has not been observed yet." + }, + "0112/astro-ph0112045_arXiv.txt": { + "abstract": "It was shown previously from the redshifts and positions of the compact, high-redshift objects near the Seyfert galaxy NGC 1068 that they appear to have been ejected from the center of the galaxy in four similarly structured triplets. In this \\em local \\em scenario, they lie at the distance of NGC 1068, a distance much closer than a cosmological interpretation of their redshifts would imply. A large portion of their measured redshifts would then be intrinsic and it was found that this intrinsic component decreases with increasing distance from the galaxy. Here some of the consequences of assuming such a $local$ model for QSOs are examined. As has been found in several similar cases, the luminosity of the objects increases systematically with the decrease in redshift. The luminosity change cannot be Doppler related and a model in which the luminosities and intrinsic redshifts vary with time is found to fit the data best. This $local$ scenario thus appears to require a model similar to the one suggested by Narlikar and Das in which the creation of matter is ongoing throughout the life of the Universe. In fact, the observed increase in luminosity with decreasing intrinsic redshift found here, is in reasonable agreement with their prediction. In their model, matter is created with a high intrinsic redshift in mini Big Bangs and is ejected in the form of QSOs from the centers of active galaxies. From the ages of the ejection events in NGC 1068 it is found that in a relatively short time ($10^{7} - 10^{8}$ yrs) the intrinsic redshift component in these objects disappears and their luminosity approaches that of a normal galaxy. This period, which is much shorter than a Hubble time, may then determine the approximate lifetime of a QSO, and, in this model, QSOs may be the first, short-lived stage in the life of a galaxy. Perhaps of even more interest is the result that, when QSOs are assumed to be local, their generation rate is found to be constant throughout the age of the Universe. There is no need to invoke an epoch of enhanced, high-luminosity QSO production as is required in the cosmological redshift model to explain the apparent bunching-up of high-luminosity QSOs with redshifts near z = 2. Finally, because QSO lifetimes are relatively short ($<10^{8}$ yrs), an initial event (Big Bang) is still required to explain the high-redshift galaxies whose intrinsic redshift component will have long since disappeared. The Hubble expansion is therefore still expected to apply for normal galaxies. ", + "introduction": "It has been argued \\citep{arp01a,arp01b}, and references therein) that the QSO-like objects clustered near and aligned along the minor axis of active galaxies have been ejected from these same galaxies. If the ejection process is to be studied in detail, however, the positions and redshifts of a large number of objects are required. So far, redshifts are not available for most of the compact objects near active galaxies considered elsewhere. Fortunately the region around NGC 1068 has been examined very closely and most of the QSOs detected near this galaxy now have measured redshifts making it one of the best places in which to search for a possible physical association. \\citet{bur99} has shown that 14 QSOs and BSOs lie within 50$\\arcmin$ of this galaxy and have redshifts from 0.261 to 2.018. \\begin{figure*} \\hspace{-2.5cm} \\vspace{-3.2cm} \\epsscale{2.1} \\plotone{f1.eps} \\caption{\\scriptsize{Sources plotted after removing N-S offsets equal to the displacements of their associated triplet mass centers. The letters A, B, C, and D identify the triplets. The \"r\" and \"b\" next to the paired sources indicates whether it is redshifted or blueshifted relative to the mean pair redshift. The direction of rotation of the triplet axes (heavy solid lines) is indicated at the top.\\label{fig1}}} \\end{figure*} \\begin{figure} \\hspace{-2.5cm} \\vspace{-2.0cm} \\epsscale{1.1} \\plotone{f2.eps} \\caption{\\scriptsize{Redshift plotted versus magnitude for the paired objects. The heavy solid line shows how the mean pair redshift varies with the mean pair magnitude. Pairs are joined by dashed lines.\\label{fig2}}} \\end{figure} \\begin{figure} \\hspace{-2.0cm} \\vspace{-0.5cm} \\epsscale{1.3} \\plotone{f3.eps} \\caption{\\scriptsize{(a) Plot of mean triplet magnitude vs age. (b) Plot of z$_{\\rm mean}$ vs age. The solid line assumes a linear relation. The dashed line shows qualitatively how the intrinsic redshift component might decrease if the maximum value of z$_{\\rm mean}$ = 1.25 as found in paper I.\\label{fig3}}} \\end{figure} \\begin{figure} \\hspace{-2.5cm} \\vspace{-2.8cm} \\epsscale{1.1} \\plotone{f4.eps} \\caption{\\scriptsize{Singlet (open circles) and mean pair magnitudes (filled circles) plotted as a function of ejection angle. The solid curve is the mean of the singlet and mean pair magnitudes.\\label{fig4}}} \\end{figure} Recently, \\citet{bel02} (hereafter called paper I) has shown that 12 of the compact objects clustered near NGC 1068 appear to have been ejected from the galaxy with very modest velocities (1-2$\\times10^{4}$ km s$^{-1}$) in four similarly structured triplets that differ only in their size and orientation. In that study, ejection velocities and angles gave event ages $\\leq6\\times10^{6}$ yrs, where the event age is defined as the elapsed time since the ejection event occurred. The rotation period of the central object was also determined to be $10^{7}$ yrs, in good agreement with the rotation period of the nucleus of NGC 1068 obtained using H$_{2}$ results \\citep{all01}. It was also shown that the measured redshifts are likely to be composed of several Doppler components and one non-Doppler, or intrinsic, component. Although source pairs have been reported previously \\citep{arp97a,arp97b,arp98}, reports of triplets \\citep[see][]{bur80} are less common. Here some of the consequences that result from the assumption that the QSOs near NGC 1068 have been ejected from it are examined. \\begin{deluxetable}{cccccccc} \\tabletypesize{\\scriptsize} \\tablecaption{Redshifts, magnitudes and ages of the triplets near NGC 1068. \\label{tbl-1}} \\tablewidth{0pt} \\tablehead{ \\colhead{Triplet(pair,singl.)} & \\colhead{Eject. Angle($\\arcdeg$)\\tablenotemark{a}} & \\colhead{Mean pr. z} & \\colhead{z$_{(\\rm r-b)}$} & \\colhead{Pair mag.} & \\colhead{Singl. Mag} & \\colhead{Tripl. mag} & \\colhead{Age$(10^{6})$yrs\\tablenotemark{b}} } \\startdata A(1-3,14)\\tablenotemark{c} & -122 & 0.4935 & 0.1556 & 17.1 & 18.2 & 17.46 & 6.08 \\\\ B(2-5,13) & -71 & 0.7610 & 0.1664 & 18.4 & 18.2 & 18.33 & 4.46 \\\\ C(4-6,12) & -38 & 1.1005 & 0.2150 & 18.5 & 19.3 & 18.76 & 3.42 \\\\ D(8-11,10) & -4.4 & 1.2015 & 0.371 & 18.6 & 19.9 & 19.03 & 2.81 \\\\ \\enddata \\tablenotetext{a}{Measured clockwise from l-o-s for $\\gamma = 18^{\\arcdeg}$} \\tablenotetext{b}{Assumes rotation axis tip angle $\\gamma$ = $18^{\\arcdeg}$ \\citep{bel02}} \\tablenotetext{c}{Source numbers as listed in \\citet{bur99}} \\end{deluxetable} ", + "conclusions": "Using the results found previously for the QSOs near NGC 1068 \\citep{bel02} it is shown that, in this $local$ model, changes in their magnitudes and redshifts with time during the QSO stage can be fitted best to a model similar to that proposed by \\citet{nar80b}. The compact objects near NGC 1068 (z$_{\\rm c}$ = 0.0038) are born with apparent magnitudes near 19.5 (or fainter) and a large intrinsic redshift component that decreases as their luminosity increases. After $\\sim10^{8}$ yrs this intrinsic redshift component will have largely disappeared and they are assumed to evolve into galaxies. In this scenario, the formation of QSOs (and therefore also galaxies) is \\em continuous \\em and \\em uniform \\em throughout the entire age of the Universe. These results (varying luminosity, redshift and mass) differ from the model of Narlikar and Das in that they seem to apply only during the object's birth and the first brief period thereafter (the QSO stage). Because the QSO stage is so short, an initial Big Bang still appears to be required to explain the existence of high-redshift galaxies whose intrinsic redshift component should have long ago fallen to near-zero values. It has been shown that in this model it is no longer necessary to hypothesize a period of high-luminosity QSO production near z = 2, as is required for the cosmological redshift model. In the QSO stage, the change of luminosity with intrinsic redshift and the change of mass with time are both found to agree well with the predictions of the Narlikar and Das theory. The amount of time spent in the QSO stage ($\\sim10^{8}$ yrs) compared to the life of a galaxy, results in a galaxy/QSO number ratio of $\\sim10^{2} - \\sim10^{3}$. Finally, it is suggested that, in the future, if this $local$ model is ever to be convincingly confirmed, or proven to be incorrect, it is very important to continue to measure the redshifts and magnitudes of those QSOs that appear to have been ejected from nearby active galaxies. This is especially true for those cases where a large number of objects have been reported. \\newpage" + }, + "0112/astro-ph0112209_arXiv.txt": { + "abstract": "Globular cluster systems of 17 elliptical galaxies have been studied in the Coma cluster of galaxies. Surface-brightness fluctuations have been used to determine total populations of globular clusters and specific frequency ($S_N$) has been evaluated for each individual galaxy. Enormous differences in $S_N$ between similar galaxies are found. In particular, $S_N$ results vary by an order of magnitude from galaxy to galaxy. Extreme cases are the following: a) at the lower end of the range, NGC 4673 has $S_N = 1.0 \\pm 0.4$, a surprising value for an elliptical galaxy, but typical for spiral and irregular galaxies; b) at the upper extreme, MCG +5 $-$31 $-$063 has $S_N = 13.0 \\pm 4.2$ and IC 4051 $S_N = 12.7 \\pm 3.2$, and are more likely to belong to supergiant cD galaxies than to ``normal\" elliptical galaxies. Furthermore, NGC 4874, the central supergiant cD galaxy of the Coma cluster, also exhibits a relatively high specific frequency ($S_N = 9.0 \\pm 2.2$). The other galaxies studied have $S_N$ in the range [2, 7], the mean value being $S_N = 5.1$. No single scenario seems to account for the observed specific frequencies, so the history of each galaxy must be deduced individually by suitably combining the different models ({\\it in situ}, mergers, and accretions). The possibility that Coma is formed by several subgroups is also considered. If only the galaxies of the main subgroup defined by \\citet{GM01} are used, a trend in $S_N$ arises in the sense of $S_N$ being bigger in higher density regions. This result needs further confirmation. ", + "introduction": "Globular clusters (GCs) are thought to be among the oldest objects in the Universe; so, although they represent only a small fraction of the total luminosity of the galaxy hosting them, they provide useful information about the galaxy formation process and its initial evolutionary stages. By studying globular cluster systems (GCS), the memory of the host galaxy is probed. A fundamental, widely used parameter describing a GCS is the specific frequency, $S_N$, introduced by \\citet{HB81}. $S_N$ is the number of GCs normalized to an absolute magnitude for the parent galaxy of $M_V^{\\rm TOT}=-15$ mag: \\begin{equation} S_N=N_{\\rm GC}^{\\rm tot} 10^{0.4(M_V^{\\rm TOT} + 15)}, \\end{equation} where $N_{\\rm tot}$ is the total number of GCs in the galaxy. $S_N$ gives the number of GCs per unit luminosity of the host galaxy. Spiral and irregular galaxies have typical values of $S_N \\le 1$ \\citep{BH82}. Dwarf \\citep{D96,M98}, elliptical \\citep{HB81}, and lenticular galaxies \\citep{KW98,C99} have $S_N$ between 2 and 6, where larger values seem to correspond to galaxies located in high-density environments \\citep{H91}. Giant cDs located near the centers of rich clusters of galaxies have the largest specific frequency, $S_N \\sim 10 - 20$ (Harris \\& van den Berg 1981; Harris, Pritchet, \\& McClure 1995; Blakeslee \\& Tonry 1995; Bridges et al. 1996), with some exceptions \\citep{K96}. The origin of the observed values of $S_N$ is not clear. A basic question is whether $S_N$ depends on the properties of the host galaxy and/or the environment. If only elliptical galaxies are considered, environmental density seems to play a fundamental role in $S_N$ \\citep{H91}. These results, which are based on the compilation of several studies of elliptical galaxies closer than the Virgo cluster, suggest that the formation of GCs has been two or three times more efficient in rich environments than in poor ones. On the other hand, the luminosity of the host galaxy also has an influence on the properties of the GCS (Harris, Harris, \\& McLaughlin 1998). \\citet{F82} and \\citet{M87} have suggested that in clusters of galaxies a population of intergalactic GCs (IGCs) should exist. A galaxy could enlarge its GC population via gravitational capture of IGCs. If they are distributed according to the gravitational potential of the whole galaxy cluster, then the influence of the environment on $S_N$ is justified \\citep{W87,W95}. On the other hand, \\citet{B97} suggested that $S_N$ could be strongly influenced by gas removal by tidal disruption during the formation of the cluster. This idea was developed further by Blakeslee, Tonry, \\& Metzger (1997). The star formation would be inhibited as a result of the gas loss, increasing $S_N$. Moreover, \\citet{H98} argued that explaining the observed $S_N$ in M87 via gravitational capture of IGCs, requires assuming an unrealistic spatial distribution for intergalactic GCs. The origin of the GCSs themselves in giant elliptical galaxies is still not understood. One of the main problems in understanding this phenomenon is the lack of data. The difficulty in observing distant GCSs is due to the low resolution of available telescopes and to the low brightness of the individual GCs, which are too faint to be detected even at moderate distances. At $\\sim 100$ Mpc (the distance of the Coma galaxy cluster), only the brightest GCs are visible from ground-based observations (\\citet{H87} and \\citet{TV87} performed the first studies of GCSs in Coma); high-resolution telescopes (such as the {\\it Hubble Space Telescope} ({\\it HST\\/}), see for example \\citet{H00}, \\citet{K00}) and long exposure times are required to detect them directly. But the surface-brightness fluctuations (SBF) technique, which we use in this paper, allows us to detect GCSs in elliptical galaxies at distances of the order of 100 Mpc, using small to intermediate-size, ground-based telescopes and reasonable exposure times \\citep{BT95}. The SBF technique was introduced by \\citet{TS88} with the aim of measuring distances. Comparing the fluctuation signals produced by the stellar population of a galaxy with SBF measurements of nearby galaxies which have externally calibrated distances, accurate estimates of distances can be obtained up to $\\sim$ 40 Mpc \\citep{T00,T01}. For smaller distances, undetected GCs contribute as a perturbation of the total fluctuation signal of the galaxy. But in the case of more distant galaxies the fluctuation signal produced by the GCs, which are mostly unresolved, dominates and the contribution of stars becomes negligible (Wing et al. 1995; Blakeslee \\& Tonry 1995; Blakeslee, Tonry, \\& Metzger 1997; Blakeslee 1999). In this paper, we use the SBF technique to evaluate the GCS properties of elliptical galaxies in the Coma cluster. Coma is a good laboratory for studying GCSs, because it is a rich cluster and its ellipticals have a wide range of luminosities and environmental densities. After describing the SBF technique and its use in obtaining GCS properties (\\S \\ref{SBF}), we present a test for evaluating the capacity of the technique (\\S \\ref{test}), which we then apply to the Coma cluster galaxies (\\S \\ref{result}). ", + "conclusions": "\\label{conclusions} In the following, we discuss first the scenarios for elliptical-galaxy formation in the light of the obtained results, and second, the possibility that a merging process is at work in Coma. \\subsection{Elliptical Galaxy Formation Scenarios} We have measured $S_N$ in 17 elliptical galaxies located in Coma. In Table \\ref{sn} we present a summary of the obtained results. For each galaxy we show the value of $M_V^{\\rm TOT}$ (obtained from the RC3 value and adopting a distance modulus for Coma of $34.89 \\pm 0.20$), the distance to the central galaxy, NGC 4874, the obtained total population of GCs, and the result for $S_N$. The uncertainty in the distance has not been included in the results, it produces an additional error of 18\\% in $S_N$. These results reveal enormous differences in $S_N$ among similar galaxies. In particular, $S_N$ varies by an order of magnitude from galaxy to galaxy. Extreme cases are: a) at the lower end, NGC 4673 has $S_N = 1.0 \\pm 0.4$, this value being typical of spiral or irregular galaxies, but surprising for an elliptical galaxy; b) on the other hand, MCG +5 $-$31 $-$063 has $S_N = 13.0 \\pm 4.2$ and IC4051 has $S_N = 12.7 \\pm 3.2$, which are similar to the values found in supergiant cD galaxies, but not in ``normal\" elliptical galaxies. \\citet{BTM97} also performed an SBF analysis and reported $S_N$ results for NGC 4874, NGC 4889 and NGC 4839. They obtained $S_N=9.3 \\pm 2.0$ for NGC4874; $5.7 \\pm 1.3$ for NGC 4889, and $4.6 \\pm 1.5$ for NGC 4839. These results are compatible (within the error bars) with ours. We obtained for these galaxies $S_N=9.0 \\pm 2.2$, $S_N=4.0 \\pm 1.2$ and $S_N=7.0 \\pm 1.9$ respectively. Perhaps there is a small difference in the case of NGC 4839. Note that results reported by \\citet{BTM97} are ``$metric-S_N$\"; i.e., $S_N$ calculated within a radius of 40 kpc. If GCs and halo light follow the same radial distribution, $S_N$ does not change with radius and the ``$metric-S_N$\" would be identical to the global $S_N$. The results for NGC 4839 probably indicate that NGC 4839 GCS is more extended than its halo. Indeed, NGC 4839 has the most extended GCS of our sample (see Fig. \\ref{radiales.grafi}). \\citet{WH00} studied {\\it HST} images of IC 4051 and proved that a central location in a rich cluster environment is not required to form a high population of GCs. They obtained a $S_N$ equal to $11 \\pm 2$. In our work, IC 4051, a ``normal\" elliptical near the cluster core, has a high $S_N$ of 12.7 $\\pm$ 3.2. Furthermore, NGC 4874, the central supergiant cD galaxy in Coma, also exhibits a relatively high specific frequency, $S_N = 9.0 \\pm 2.2$. The remaining galaxies studied have $S_N$ in the range [2, 7], the mean value being $S_N = 5.1$. Why do IC 4051 and NGC 4673, galaxies with similar absolute magnitudes, have differences of a factor of twelve in $S_N$? Perhaps this is due to differences in the environment. But then, why does MCG +5 $-$31 $-$063 have $S_N$ five times bigger than IC 4041, if both galaxies are located on the border of the Coma cluster core and have the same absolute magnitudes?. In order to study possible relations between $S_N$ and environment, $S_N$ versus the distance $R$ to the central Coma galaxy, NGC 4874 is plotted in Fig. \\ref{snr}. No clear trend is found in this plot, which suggests that $S_N$ does not depend significantly on the environment in Coma. On the other hand, Fig. \\ref{snm} shows $S_N$ versus $M_V^{\\rm TOT}$ of each galaxy. No relation is found between $S_N$ and $M_V^{\\rm TOT}$. The figure is completely dominated by the dispersion of the points, which is greater than the error bars and must therefore be real. Formation scenarios for giant ellipticals tend to fall into three basic classes of models: a) in situ models, in which the galaxy condenses by dissipative collapse of gas clouds in one or more major bursts, b) mergers of gas-rich systems, probably disk-type galaxies, and c) accretion of smaller satellites. Various combinations of these extremes are also possible. In situ models predict a correlation between galaxy properties and the GCS. If we assume a universal efficiency of GC formation per unit total initial galaxy mass \\citep{Mc99}, \\begin{equation} \\epsilon = \\frac{M_{\\rm GC}}{M_* + M_{\\rm gas}} \\end{equation} where $\\epsilon$ is the efficiency parameter, $M_*$ is the mass of the visible stellar component of the galaxy, $M_{\\rm gas}$ is the mass of gas in or around the galactic halo, and $M_{\\rm GC}$ is the mass of the GCS, then $S_N$ must be related with the luminosity of the host galaxy according to: \\begin{equation} S_N \\sim \\epsilon (1+\\frac{M_{\\rm gas}}{M_*})L_{\\rm gal}^{0.3}. \\end{equation} The term $L_{\\rm gal}^{0.3}$ accounts for the systematic increase in mass-to-light ratio with galaxy luminosity. \\citet{Mc99} predicts the general trend of $S_N$ with luminosity. The only unknown parameter affecting $S_N$ would be the ratio of gas to stellar mass. In this paper no evidence is found for a relation between galactic luminosity and $S_N$. On the other hand, if a range of possible GC formation efficiencies is allowed, the merger model can account for the specific frequency range observed for elliptical galaxies \\citep{AZ92}. If the merger model is to explain the $high-S_N$ phenomenom, GC formation must be more efficient in the $high-S_N$ galaxies. \\citet{F82} and \\citet{M87} have suggested that in clusters of galaxies a population of intergalactic GCs (IGCs) should exist. A galaxy could enlarge its GC population via gravitational capture of IGCs. If they are distributed according to the gravitational potential of the whole galaxy cluster, a relation between $S_N$ and environment is predicted \\citep{W87}. The current data set shows no indication supporting this prediction. The fact that no single scenario seems to account for the observed specific frequencies, indicates the history of each galaxy should be deduced individually by suitably combining the models mentioned above \\citep{WH00}. To this aim, it becomes necessary to extend the observational information on each galaxy. Besides the specific frequency and the radial distribution of the GCS, it is necessary to know the details of the subpopulations (when these exist) and the kinematics of the GCS. In this way, detailed {\\it HST} observations of the extreme cases mentioned above (NGC 4673 and MCG +5 $-$31 $-$063) could be a good starting point because these peculiar galaxies can show characteristics found nowhere else, and so offer valuable information in testing the different scenarios. \\subsection{Subgroups and Merging in Coma} \\citet{GM01} have recently discovered the existence of three subgroups of galaxies in Coma, one of them associated with the cD galaxy NGC 4874 and the other two with NGC 4889 and NGC 4839. They conclude that the non-stationarity of the dynamical processes at work in the Coma core is due to the merging of small-scale groups of galaxies. In this context, each subgroup formed separately and then the merger between the different groups took place. If this scenario is valid, is there any relation between $S_N$ and environment and/or $M_V^{\\rm TOT}$ inside each subgroup? In order to analyze this question, we restricted our figures to galaxies belonging to \\citet{GM01} subgroup 2 and studied in this paper: NGC 4874, IC 4012, IC 4041, IC 3976, and IC 3959. In Fig. \\ref{snmv2} $S_N$ versus host galaxy magnitude is plotted, while Fig. \\ref{snr2} shows $S_N$ versus the distance to the galaxy NGC 4874, which is very close to the center of subgroup 2. No relation between $S_N$ and $M_V^{\\rm TOT}$ is found from Fig \\ref{snmv2}. But there is an apparent trend in Fig. \\ref{snr2}: $S_N$ is bigger in high density environments. Is this trend real or is it an artifact of the low number of galaxies considered?. If this result is confirmed, this will be a strong argument in favor of the IGCs model, so the next step in this work will be to enlarge the number of galaxies of the study." + }, + "0112/astro-ph0112515_arXiv.txt": { + "abstract": "{ A comprehensive temporal analysis has been performed on the 319 brightest GRBs with $T_{90}$$>$$2\\,$s from the BATSE current catalog. The GRBs were denoised using wavelets and subjected to an automatic pulse selection algorithm as an objective way of identifying pulses and quantifying the effects of neighbouring pulses. The number of statistically significant pulses selected from the sample was greater than 3000. The rise times, fall times, full-widths at half maximum (FWHM), pulse amplitudes and pulse areas were measured and the frequency distributions are presented here. All are consistent with lognormal distributions provided the pulses are well separated. The distribution of time intervals between pulses is not random but compatible with a lognormal distribution when allowance was made for the 64 ms time resolution and a small excess (5\\%) of long duration intervals that is often referred to as a Pareto-L\\'{e}vy tail. The time intervals between pulses are most important because they may be an almost direct measure of the activity in the central engine. Lognormal distributions of time intervals also occur in pulsars and SGR sources and therefore provide indirect evidence that the time intervals between pulses in GRBs are also generated by rotation powered systems with super-strong magnetic fields. \\\\ [\\parsep] \\indent A range of correlations are presented on pulse and burst properties. The rise and fall times, FWHM and area of the pulses are highly correlated with each other. The pulse amplitudes are anticorrelated with the FWHM. The time intervals between pulses and pulse amplitudes of neighbouring pulses are correlated with each other. It was also found that the number of pulses, N, in GRBs is strongly correlated with the fluence and duration and that can explain the well known correlation between duration and fluence. The GRBs were sorted into three categories based on N i.e. 3$\\leq$N$\\leq$12, 13$\\leq$N$\\leq$24 and N$\\geq$25. The properties of pulses before and after the strongest pulse were compared for three categories of bursts. No major differences were found between the distributions of the pulse properties before and after the strongest pulse in the GRB. However there is a strong trend for pulses to have slower rise times and faster fall times in the first half of the burst and this pattern is strongest for category N. This analysis revealed that the GRBs with large numbers of pulses have narrower and faster pulses and also larger fluences, longer durations and higher hardness ratios than the GRBs with smaller numbers of pulses. These results may be explained by either homogeneous or inhomogeneous jet models of GRBs. The GRBs with larger number of pulses are closer to the axis if $\\Gamma$ varies with the opening angle of the jet and the imprint of the jet is preserved in the pulse structure of the burst. The distribution of the number of pulses per GRB broadly reflects the beaming by the jet. ", + "introduction": "Much of the recent progress in the study of gamma-ray bursts (GRBs) results from the detection of bursts with good location accuracy by BeppoSAX that enabled the detection of counterparts at other wavelengths. The subsequent redshift determination of bursts have established that these bursts are at cosmological distances \\citep{cfpc:1997,vanpara:1997}. GRBs seem to be connected to massive stars and become powerful probes of the star formation history of the universe \\citep{lamb:2000,han:2000,berger:2001}. However not many redshifts are known and there is still much work to be done to determine the mechanisms that produce these enigmatic events. The most plausible GRB progenitors are expected to be a newly formed black hole (BH) surrounded by a temporary accretion disk \\citep{rees:1999,mes:2001,castro:2001}. The most popular models include the merger of a neutron star (NS) and a NS \\citep{elp:1989,ruffjan:1999}, NS and a BH \\citep{pacz:1991}, BH white dwarf merger \\citep{fryer:1999} and models of failed supernovae or collapsars \\citep{macfad:1999,pacy:1998}. An important exception is the model in which the GRB energy is provided by a newly formed neutron star \\citep{usov:1992,thomp:1994}. Various explanations have been put forward for the complicated structure of the light curves. These range from internal shocks, caused by variations in the velocity of the outflow \\citep{reemes:1994,piran:1999}, to external shocks, caused by interactions with an external medium \\citep{meszar:1993,derm:1999}. In the internal shock model the instabilities in the wind leads to shocks which convert a fraction of the bulk kinetic energy to internal energy remote from the central engine. A turbulent magnetic field then accelerates electrons which radiate by synchrotron emission and inverse Compton scattering, generating the GRB. Many of the observed features in bursts can be reproduced in the internal shock models of GRBs \\citep{sapi:1997,kps:1997,daimoc:1998,pansm:1999,downes:2001}. A variety of analytical techniques has been applied to the temporal and spectral profiles of GRBs which place constraints on the observed distributions which models must satisfy. The impressive results from these studies include (1) hard to soft evolution \\citep{golens:1983,borgon:2001}; (2) the duration-hardness anticorrelation \\citep{kmf:1993}; (3) the temporal asymmetry of pulses in GRBs \\citep{nnw:1993,lp:1996}; (4) a bimodal duration distribution of GRBs consistent with two lognormal distributions \\citep{kmf:1993,mhlm:1994}; (5) the discovery of two different types of pulses in GRBs\\citep{ppb:1997}; (6) a correlation between E$_{\\rm peak}$ and intensity\\citep{mpp:1995}; (7) energy dependence of the pulse duration \\citep{nnb:1996}; (8) a relationship between the pulse peak energy, E$_{\\rm peak}$, and the photon fluence \\citep{lika:1996,crider:1999}; (9) lognormal pulse shapes and time intervals between pulses in long \\citep{mhlm:1994,hmq:1998} and short GRBs \\citep{sheila:2001}; (10) spectra well fit with a Band function \\citep{band:1993}; (11) spectral hardening before a count rate increase \\citep{bhat:1994}; (12) an x-ray excess in GRB spectra \\citep{stroh:1998}; (13) a correlation between complexity and brightness \\citep{stern:1999} and (14) the unique properties of the pulses and power law relationships between the pulse properties and durations of GRBs \\citep{smcb:2002} While GRBs display hard to soft spectral evolution, there is remarkable constancy of the pulses in GRBs throughout the burst \\citep{ramfen:2000,qhm:1999}. The temporal and spectral properties of a few GRBs with known redshift have yielded two important results to suggest that GRB properties may be related to their luminosities. Ramiriz-Ruiz and Fenimore (1999) have shown that more rapidly variable bursts have higher absolute luminosities. \\citet{nmb:2000} have found an anticorrelation between the time delay in the arrival times of hard and soft photons in pulses and the luminosity of the GRB. The light curves of GRBs are irregular and complex. Statistical studies are necessary to characterise their properties and hence to identify the physical properties of the emission mechanism. The statistical methods used for temporal studies can be broadly divided into four categories: (1) fits to individual pulses in the GRB using a number of pulse shape parameters \\citep{nnb:1996,lee:2000,lbp:2000}; (2) a non-parametric approach to pulse shapes in GRBs \\citep{mhlm:1994,hmq:1998,ymr:1995,qhm:1999}; (3) the average statistical properties of GRBs using a peak-aligned profile \\citep{ss:1996}; and (4) the average power spectral density of GRBs \\citep{belli:1992,belss:2000,changyi:2000}. One of the first studies \\citep{mhlm:1994} revealed that lognormal distributions can adequately describe the properties of GRBs. Subsequent studies \\citep{lifen:1996,hmq:1998,qhm:1999} have confirmed the applicability of lognormal distributions in accounting for the wide range in the observed properties of pulses in GRBs. This result is not surprising because lognormal distributions arise from the product of probabilities of a combination of independent events and such conditions apply to the pulse generation process in GRBs. In a different approach \\citep{belss:2000} used Fourier analysis to study the power spectral density of long GRBs. This approach revealed that the diversity of GRBs is due to random realisations of the same process which is self-similar over a range of time scales \\citep{ss:1996}. The slope of the PSD was -5/3 suggesting that GRBs are related to fully developed turbulence. The two different approaches are quite similar because the lognormal approach has been used to describe fully developed turbulence \\citep{amm:1999}. The work presented here expands on the earlier analysis \\citep{quilligan:2000} and provides new insight into the mechanism which generates GRBs. The aim is to provide a comprehensive description and understanding of the pulse properties in GRBs and combine it with other studies of the spectral properties. The wavelet analysis and the pulse selection algorithm are described in Sect. 2. The method for comparing the properties of the pulses before and after the strongest pulse in the GRB is also described in Sect. 2. The results are presented in Sect. 3, and discussed in Sect. 4. The conclusions are presented in Sect. 5. ", + "conclusions": "The properties of the brightest 319 GRBs in the BATSE current catalogue have been analysed.The automatic pulse selection process detected more than 3300 pulses. The distributions of pulse rise and fall times, FWHM, areas, amplitudes and time intervals between pulses are reasonably consistent with the lognormal distribution. GRB pulse profiles can be elegantly described by a small number of parameters that may be very useful for simulations. The lognormal distribution depends on the product of probabilities arising from a combination of independent events and these conditions must therefore apply to the generation of the temporal and spectral properties of GRB pulses. A wide range of burst parameters and also pulse parameters were correlated and the results follow the trend expected from the internal shock model. The pulse amplitude is strongly anticorrelated with the other pulse timing parameters. The time intervals between pulses and pulse amplitudes are correlated with each other. A comprehensive analysis has been performed between the first half and second half of GRBs in three categories defined in terms of N. No major differences were found between the distribution of pulse properties between the first half and second half of the GRBs. There is a strong tendency for pulses to have slower rise times and faster fall times in the first half of the burst. This trend is stronger in GRBs with small numbers of pulses. The pulse timing parameters and time intervals all decrease with increase in N. These results seem to be compatible with jet models with either a $\\Gamma$ that varies with the opening angle or is constant and varies with the mass. If $\\Gamma$ varies with the opening angle of the jet, the GRBs with higher values of $\\Gamma$ and greater variability are observed close to the axis of the jet while GRBs with smaller number of pulses and less variability are observed at larger angles from the jet. Jets with values of $\\Gamma$ that vary with angle or with mass may explain the luminosity-variability correlation and the luminosity-energy lag correlation in GRBs with known redshift. This study of the number of pulses in GRBs and their time structure provides strong evidence for rotation powered systems with intense magnetic fields and the added complexity of a jet. These results can be well interpreted by internal shocks in the framework of theoretical models for the formation of black holes and subsequent jet formation." + }, + "0112/astro-ph0112386_arXiv.txt": { + "abstract": "Powerful quasars and radio galaxies are injecting large amounts of energy in the form of radio plasma into the inter-galactic medium (IGM). Once this nonthermal component of the IGM has radiatively cooled the remaining radio emission is difficult to detect. Two scenarios in which the fossil radio plasma can be detected and thus be used to probe the IGM are discussed: a) re-illumination of the radio emission due to the compression in large-scale shock waves, and b) inverse Compton scattered radiation of the cosmic microwave background (CMB), cosmic radio background (CRB), and the internal very low frequency synchrotron emission of the still relativistic low energy electron population. We present 3-D magneto-hydrodynamical simulations of scenario a) and compare them to existing observations. Finally, we discuss the feasibility of the detection of process b) with upcoming instruments. ", + "introduction": "\\begin{figure}[t] \\begin{center} \\includegraphics[width=\\textwidth]{ensslint1.eps} \\end{center} \\caption[]{Shock passage of a hot, magnetized bubble (a radio ghost) through a shock wave. The flow goes from the left to the right. The evolution of the mid plane gas density is displayed (white is dense, black is dilute gas)} \\label{eps1} \\end{figure} The jets of powerful radio galaxies inflate large cavities in the IGM that are filled with relativistic particles and magnetic fields. Synchrotron emission at radio frequencies reveals the presence of electrons with GeV energies. These electrons have radiative lifetimes of the order of 100 Myr before their observable radio emission extinguishes due to radiative energy losses. The remnants of radio galaxies and quasars are called \\lq fossil radio plasma' or a \\lq radio ghosts' \\cite{1999dtrp.conf..275E}. Their existence as a separate component of the IGM is supported by the detections of cavities in the X-ray emitting galaxy cluster gas \\cite[and others]{1993MNRAS.264L..25B,1994MNRAS.270..173C,1998ApJ...496..728H,2000ApJ...534L.135M,2000MNRAS.318L..65F,2001ApJ...547L.107F,fabian2001moriond,McNamara2000Paris,heinz2001,schindler2001}. In many cases associated radio emission and in a few cases a lack of such emission was found, as expected for aging bubbles of radio plasma. Such bubbles should be very buoyant and therefore rise in the atmosphere of a galaxy cluster. It is not clear yet if they break into pieces during their ascent and thereby are slowed down. Another possibility is that they are able to ascend up to the accretion shock of a galaxy cluster, where their further rise will be prohibited by the infalling gas of the accretion onto the cluster. ", + "conclusions": "" + }, + "0112/astro-ph0112453_arXiv.txt": { + "abstract": "This paper will give a short description of a `cookbook' for the XMM-Newton data reduction software XMMSAS. This `cookbook' has been developed at the Max-Planck-Institut f\\\"ur extraterrestrische Physik (MPE), Garching. The task of this `cookbook' is to describe the necessary XMMSAS tasks and show examples to make it easy for new users of the XMMSAS to understand the steps which have to perform to reduce their XMM X-ray data. ", + "introduction": "When working with X-ray data each X-ray mission has its own software package to reduce the data due to the specific requirements of the detectors onboard. For the XMM-Newton mission the reduction software package XMMSAS has been developed. The main purpose of the XMMSAS is to allow users to take the Observational Data File (ODF) and create event files and at a later stage scientific files like images, spectra and lightcurves. Each XMMSAS task has a manual either in HTML format or as a postscript file (please see at \\newline {\\it xmm.vilspa.esa.es/sas/current/doc/packages.All.html} \\newline to get the whole list of available XMMSAS tasks). The manual gives a description of the tasks and lists all available task parameters. However, this is often confusing to new users of the XMMSAS in order to see which tasks are needed and how they have to be applied to their data. For this reason at MPE the idea was born to design a web-based cookbook that explains users how to work with their XMM-Newton observational data files (ODF) in order to create final scientific files like spectra and light curves. At first this cookbook was only written for users at MPE, but it has been now designed and make accessible to users from outside the MPE. ", + "conclusions": "" + }, + "0112/astro-ph0112179_arXiv.txt": { + "abstract": "Iron K$\\alpha$ emission from photoionized and optically thick material is observed in a variety of astrophysical environments including X-ray binaries, active galactic nuclei, and possibly gamma-ray bursts. This paper presents calculations showing how the equivalent width (EW) of the \\fe\\ line depends on the iron abundance of the illuminated gas and its ionization state -- two variables subject to significant cosmic scatter. Reflection spectra from a constant density slab which is illuminated with a power-law spectrum with photon-index $\\Gamma$ are computed using the code of Ross \\& Fabian. When the \\fe\\ EW is measured from the reflection spectra alone, we find that it can reach values greater than 6~\\kev\\ if the Fe abundance is about 10$\\times$ solar and the illuminated gas is neutral. EWs of about 1~\\kev\\ are obtained when the gas is ionized. In contrast, when the EW is measured from the incident+reflected spectrum, the largest EWs are $\\sim 800$~\\kev\\ and are found when the gas is ionized. When $\\Gamma$ is increased, the \\fe\\ line generally weakens, but significant emission can persist to larger ionization parameters. The iron abundance has its greatest impact on the EW when it is less than 5$\\times$ solar. When the abundance is further increased, the line strengthens only marginally. Therefore, we conclude that \\fe\\ lines with EWs much greater than 800~eV are unlikely to be produced by gas with a supersolar Fe abundance. These results should be useful in interpreting \\fe\\ emission whenever it arises from optically thick fluorescence. ", + "introduction": "\\label{sect:intro} There are a number of different astrophysical environments where optically thick material may be irradiated by X-rays. If the incident spectrum has significant flux above 7.1~\\kev, then, as a result of its relatively high cosmic abundance and large fluorescent yield, the iron~K$\\alpha$ line is predicted to be a significant feature in the resulting emission from the illuminated surface \\citep*[e.g.,][]{gf91,mpp91}. The presence of this line has been discussed for a number of different reprocessors including the solar surface \\citep{bai79}, the surface of the companion star in X-ray binaries \\citep*{bst74,b78,p79}, and the surface of accreting magnetic white dwarfs \\citep*[e.g.,][]{sfr84,dob95,vkh96}. Most famously, \\fe\\ emission has been observed in the X-ray spectra of many Active Galactic Nuclei (AGN) and black hole candidates \\citep*{np94,eb96}. Here, the irradiated surface is likely the accretion disc feeding the black hole, and some of the observed \\fe\\ lines have been observed to be broadened in a manner consistent with material orbiting in a relativistic gravitational potential \\citep*{tan95,nan97,fab00}. Finally, tentative detections of line emission, most likely from \\fe, have been found in recent observations of the X-ray afterglows of some $\\gamma$-ray bursts \\citep*[e.g.,][]{pir99,pir00,ant00}. If this is confirmed then it implies that optically thick material may be in the vicinity --- a possible constraint on the environment and progenitor of the burst \\citep*{ree00,vie01,brr01}. As one of the few physical diagnostics in the X-ray spectra of these systems, it is important to understand how the \\fe\\ emission line depends on the properties of the irradiated gas. Previous studies have examined the influence of changes in metal abundance \\citep*{mfr97}, the inclination angle of the reflector \\citep*{ghm94}, ionization state \\citep*{mfr93,mfr96}, and the density distribution of the gas \\citep*{ros99,nay00,br01} on the line emission. This paper generalizes the results of Matt and colleagues by presenting how the strength of the \\fe\\ line (presented as an equivalent width [EW]) depends on both the Fe abundance and the ionization state of the gas. These two variables have the greatest impact on the EW, and are likely to vary widely over different astrophysical environments. We anticipate that these results will be helpful in interpreting \\fe\\ detections whenever it arises from an optically thick surface. The following section details the calculations that were performed, and then Section~\\ref{sect:res} presents the results as contour plots of EW. A brief discussion concludes the paper in Section~\\ref{sect:discuss}. ", + "conclusions": "\\label{sect:discuss} The results presented in the previous section showed that increasing the iron abundance would increase the \\fe\\ EW, but its influence seemed to weaken as it grew. It is interesting to investigate this point since we have only considered abundances up to 10.1$\\times$ solar. Is it possible to obtain an arbitrarily large EW solely by increasing the Fe abundance? The answer to this question is found in Figure~\\ref{fig:gamma2-perion}. \\begin{figure} \\centerline{ \\includegraphics[width=0.56\\textwidth]{gamma2.0_perion_xlog.ps} } \\caption{Contours of (\\fe\\ EW/Fe abundance) in the $\\log \\xi$-Fe abundance plane. The EWs were calculated using the total incident+reflected spectra and we only show the case of $\\Gamma=2.0$. The contour lines are at 50, 100, 200, 300, 500, 700~eV. The effect of increasing the iron abundance is greatest when it is less than 3$\\times$ solar and when the slab is ionized.} \\label{fig:gamma2-perion} \\end{figure} This plot shows contours of \\fe\\ EW normalized by the Fe abundance for the case of $\\Gamma=2.0$ and a reflection fraction of one. In this case the contours show a maximum at low Fe abundance implying that the most ``power'' out of an Fe atom or ion occurs when its abundance is around solar. Increasing the amount of iron in the gas does permit the number of \\fe\\ photons emitted to increase, but it also strengthens the line destruction mechanisms of photoabsorption and scattering. Therefore, for a given flux of ionizing photons, competing atomic processes in the gas eventually decreases the efficiency of iron fluorescence. This is seen particularly when the gas is neutral: the normalized \\fe\\ EW falls slowly over a wide range in Fe abundance because K$\\alpha$ photons are destroyed by L shell absorption. At higher values of the ionization parameter, there is a more rapid drop in the normalized \\fe\\ EW. As is illustrated explicitly in Figure~\\ref{fig:line-evolution}, \\begin{figure} \\includegraphics[width=0.50\\textwidth]{lines_with_fe.eps} \\caption{Explicit examples of how the \\fe\\ line changes with Fe abundance. Total incident+reflection spectra are plotted for the three different values of $\\Gamma$ and three representative values of $\\xi$. Reflection from a slab with a 1.1, 5.1 and 10.1$\\times$ solar Fe abundance is plotted as the solid, short-dashed and long-dashed lines, respectively.} \\label{fig:line-evolution} \\end{figure} an increase in the iron abundance from 1.1 to 5.1$\\times$ solar causes a fairly big jump in the line strength, but then only a very small change from 5.1 to 10.1$\\times$ solar. In this case, only K shell absorption is important, but due to the finite number of ionizing photons, the K edge will saturate when the Fe abundance is about solar (from Fig.~\\ref{fig:gamma2-perion}). Increasing the Fe abundance beyond this point will not greatly increase the line emission because there are few photons available to ionize the newly introduced Fe atoms. Of course, increasing $\\xi$ allows more K shell ionizations to occur, and the normalized EW becomes more uniform with Fe abundance. Interestingly, the H-like \\fe\\ line at 6.97~\\kev\\ which is suppressed due to resonance scattering significantly strengthens as the abundance is increased. These results illustrate that above an Fe abundance of $\\sim$ 5$\\times$ solar there is only a very slow increase in the \\fe\\ EW. Therefore, it seems unlikely that the EW of the line can be significantly increased by considering iron abundances larger than 10$\\times$ solar." + }, + "0112/astro-ph0112122_arXiv.txt": { + "abstract": "Recently much work in studying Gamma-Ray Burst (GRB) has been devoted to revealing the nature of outburst mechanism and studies of GRB afterglows. These issues have also been closely followed by the quest for identifying GRB progenitors. Several types of progenitors have been proposed for GRBs: the most promising objects seem to be collapsars, compact object binaries, mergers of compact objects with helium cores of evolved stars in common envelope episodes, and also the recently discussed connection of GRBs with supernovae. In this paper we consider the binary star progenitors of GRBs: white dwarf neutron star binaries (WD-NS), white dwarf black hole binaries(WD-BH), helium core neutron star mergers (He-NS), helium core black hole mergers (He-BH), double neutron stars (NS-NS) and neutron star black hole binaries (NS-BH). Using population synthesis methods we calculate merger rates of these binary progenitors and we compare them to the observed BATSE GRB rate. For the binaries considered, we also calculate the distribution of merger sites around host galaxies and compare them to the observed locations of GRB afterglows with respect to their hosts. We find that the rates of binary GRB progenitors in our standard model are lower than the observed GRB rates if GRBs are highly collimated. However, the uncertainty in the population synthesis results is too large to make this a firm conclusion. Although some observational signatures seem to point to collapsars as progenitors of long GRBs, we find that mergers of WD-NS, He-NS, He-BH, and NS-NS systems also trace the star formation regions of their host galaxies, as it is observed for long GRBs. We also speculate about possible progenitors of short-duration GRBs. For these, the most likely candidates are still mergers of compact objects. We find that the locations NS-NS and NS-BH mergers with respect to their hosts are significantly different. This may allow to distinguish between these two progenitor models, once current and near future missions, such as HETE-II or SWIFT, measure the locations of short GRBs. ", + "introduction": "The last decade brought a great breakthrough in gamma-ray burst studies. The BATSE detectors on GRO have shown that GRBs are distributed isotropically on the sky, and that their brightness distribution is not consistent with a uniform source distribution in Euclidean space (Paciesas et al. 1999). Observations of GRB afterglows in X-ray, optical and radio wavelength domains (Costa et al.\\ 1997; Groot et al. \\ 1997b) led to identification of GRB host galaxies (Groot et al. \\ 1997a) and measurements of their redshifts. This has solved the long standing problem of their distance scale. While we learned that GRBs come from cosmological distances, there are still two major difficulties in understanding this phenomenon. First, we do not fully understand the physics of the outburst. Although several models have been proposed, they all have yet to meet some severe constraints imposed by observations (i.e. releasing energies of $10^{51}$--$10^{54}$\\ ergs in timescales as short as $10^{-2}\\,$s in the case of some GRBs). Second we do not know what are the astronomical objects leading to gamma-ray bursts, i.e. what are their progenitors. In recent years the black hole accretion disk model for GRBs has been given much attention (Fryer, Woosley \\& Hartmann 1999a; Meszaros 2000; Brown et al.\\ 2000). Progenitors leading to this model include collapsars (Woosley 1993; Paczynski 1998, MacFadyen \\& Woosley 1999) and binary mergers: helium star black hole (Fryer \\& Woosley 1998), double neutron stars (Ruffert et al.\\ 1997; Meszaros \\& Rees 1997), black hole neutron star (Lee \\& Kluzniak 1995; Kluzniak \\& Lee 1998) and black hole white dwarf systems (Fryer et al.\\ 1999b). Also recently the connection between supernovae and gamma-ray bursts received much attention (Paczynski 1999; Woosley 2000; Chevalier 2000), however, there is still no clear evidence that these two phenomena are intrinsically correlated (Graziani et al.\\ 1999). A good method of discerning among the binary progenitors is to compare theoretical predictions of their merger site distributions around host galaxies with location of observed GRBs within host galaxies. Binary population synthesis can be of great help in addressing this question. One can calculate the properties of a given binary population and then place it in a galactic gravitational potential to trace each binary until its components merge due to gravitational wave energy losses. This method has to deal however with a number of uncertainties, that are inherent in the binary population synthesis. Moreover there are uncertainties in what type and mass of a galaxy to use. The binary population synthesis method has already been applied to the study of compact object binaries in the context of GRB progenitors. However most studies have been concentrated only on double neutron stars and black hole neutron star systems. Lipunov et~al. (1995) have used their ``scenario machine'' to model the population of double neutron star and black hole neutron star binaries in a galaxy. They calculated the expected $\\log N$-$\\log S$ GRB distribution assuming that they are standard candles and compared it with the BATSE observations. Portegies-Zwart and Yungelson (1998) have considered the origin and properties of double neutron star systems and black hole neutron star binaries. They considered a few binary population synthesis models, with varying kick velocities, initial binary separations, initial mass ratios distribution and also considered cases with and without hyper accretion in the common envelope stage. They found the rates of mergers to be consistent with the GRB rate provided that GRBs are collimated to about ten degrees, and mentioned that double neutron stars may travel Mpc distances out of a Milky Way like galaxy before merging. Bloom, Sigurdsson \\& Pols (1999) considered double neutron stars as possible GRB progenitors, and calculated distributions of mergers of these binaries around galaxies with different masses, varying the average kick velocities in the code. They found that a significant fraction of double neutron stars merge outside their host galaxies. Bulik, Belczynski \\& Zbijewski (1999) considered the mergers of binaries containing neutron stars, and Belczynski, Bulik \\& Zbijewski (2000) investigated differences between the populations of black hole neutron star binaries and the double neutron stars. Belczynski et al.\\ (2000) found that black hole neutron star binaries merge closer to the hosts than the double neutron stars. Fryer et al.\\ (1999a) considered other types of binary progenitors of GRBs within the framework of the black hole accretion disk model of the GRB central engine. These were white dwarf black hole mergers, helium star black hole mergers, and collapsars in addition to the double neutron star systems and black hole neutron star binaries. They performed a thorough parameter study, and repeated the calculations with a number of modifications of their standard evolutionary model. They calculated the distribution of merger sites in the potentials of galaxies with the masses of $M_{MW}$ (Milky Way), $0.25 M_{MW}$, and $0.01 M_{MW}$, however they do not vary the galactic size with mass. Bloom, Kulkarni \\& Djorgovski (2001) presented a very detailed study of the observational offsets between observed afterglows and GRB hosts galaxies. They compare these observations with the theoretical distributions calculated with the code of Bloom et al.\\ (1999) and conclude that the so called delayed merging remnants i.e. double neutron star systems and black hole neutron star binaries are unlikely to be GRB progenitors, and argue in favor of the prompt bursters like collapsars and black hole helium star mergers. In this work we extend our previous studies (Belczynski \\& Bulik 1999; Bulik et al.\\ 1999; Belczynski et al.\\ 2000) to include four more proposed binary progenitors: compact object (black hole or neutron star) white dwarf binaries and Helium star mergers with black holes or neutron stars. We use a much improved and well tested binary population synthesis code and for consistency we also present the updated results for the two previously studied types of proposed progenitors: double neutron star and black hole neutron star systems. We calculate the properties of the ensemble of each type of the proposed GRB progenitors and find their distributions around different types of host galaxies. We compare the observed GRB distribution around host galaxies with the models. In order to verify the robustness of the results we perform a detailed parameter study and discuss the population synthesis models which are responsible for the largest differences. An additional way of telling which group of the proposed binaries might be responsible for GRBs is to predict their rates and compare them to the observed rate of GRBs. Population synthesis is a powerful tool for predicting rates of binary populations although it suffers from many uncertainties as some parameters of single and binary evolution are poorly known. Moreover, population synthesis works well in predicting the relative numbers of events, while calculation of absolute rates requires additional assumptions. However, such attempts have been made by a number of authors mentioned above. Using the population synthesis method we calculate merger rates of white dwarf neutron star, white dwarf black hole, double neutron star, neutron star black hole systems, and formation rates of Helium star black hole and neutron star mergers. We compare the BATSE detection rate of GRBs with the cosmic rates of the binary progenitors predicted in our calculations. In \\S\\,2 we describe the population synthesis code {\\em StarTrack} used to calculate properties of binary GRB progenitors, in \\S\\,3 we present the results, and finally \\S\\,4 is devoted to discussion with conclusions. ", + "conclusions": "We have presented calculation of rates and spatial distributions around host galaxies of several binary merger events, which were proposed as possible GRB progenitors. We have used the {\\em StarTrack} population synthesis code in our calculations. We have found that the rates are very sensitive to the assumed set of stellar evolutionary parameters. Using the rates alone we were not able to exclude any of the proposed binaries as GRB progenitors, since the highest rates obtained were always higher than observed BATSE GRB rate for any type of a binary. In the framework of the standard population synthesis model (model A) we find that the total rate of all the proposed binary events is roughly ten times larger than the observed GRB rate. However, we find that the spread in the rates due to uncertainties in population synthesis is large, and in some cases exceeds a factor of $\\sim 100$. This corresponds to the uncertainty in the estimate of the collimation of a factor of ten. On the other hand we note that our standard model (model A) leads to an expected collimation half--opening angle of $\\Theta \\ga 25^\\circ$. The measured collimation angles are somewhat smaller than this value, typically a few degrees (Panaitescu \\& Kumar 2001). Estimates of the GRB rates or collimation based on population synthesis alone carry large systematic errors. Distributions of binary system merger sites around galaxies may be compared to the locations of GRB optical afterglows with respect to the galaxies identified as their hosts. Most of GRBs take place inside or close to the host galaxies (e.g., Bloom et al.\\ 2001). Observed GRB hosts are small-mass galaxies, often thought to be going through vigorous star formation phase. There are no reliable GRB host mass estimates, and thus we have calculated models for a range of galaxy masses. Our standard model calculations were repeated for a number of different evolutionary models to assess the robustness of our results. We have found that the NS-BH mergers take place mainly outside of their host galaxies, and thus are inconsistent with the observed locations of GRBs around hosts. Some WD-BH binaries may merge outside the star formation regions of their host galaxies. However, the distribution of the WD-BH merger sites around their host galaxies is consistent with the observed distribution of GRB offsets from the centers of galaxies. Thus one can not reject the WD-BH mergers purely on the basis of comparison with the observed offsets. However, if one additionally requires that the mergers should take place in the proximity of star forming regions in galaxies then the WD-BH mergers can be rejected as potential GRB candidates. Merger sites of WD-NS, He-NS, He-BH, and NS-NS trace the star formation regions of the hosts, for all the cases of the host mass and size considered here, and independently of the adopted population synthesis model. We conclude that these types of binaries may be responsible at least for a part of observed GRBs. GRBs form a very nonuniform group of events, with different outburst times, very different light curves and observed energies. Thus, there is a possibility that GRBs originate in more than one type of progenitor. Locations of GRBs with respect to host galaxies has so far been measured only for long GRBs. There is a growing evidence that these GRBs are related to collapsing massive stars: collapsars. However, our results show that several types of binary system progenitors cannot be rejected purely on the basis of their merger site distribution. Additionally, if binaries were responsible for only a part of the observed GRBs, we also can not exclude them purely on the basis of their expected coalescence rates. Because of the expected short duration times, NS-NS and NS-BH mergers are the primary candidates as short burst progenitors. These two populations exhibit very different distributions of merger sites. Mergers of NS-NS systems take place predominantly within hosts, to the contrary of what was so far believed, provided that CE phases initiated by low-mass helium stars do not always lead to binary component mergers, assumption which yet have to be tested by detailed hydrodynamical calculations. On the other hand, a significant fraction of NS-BH systems merge outside of their host galaxies. At some point in the future afterglows from short GRBs will be observed and their locations with respect to host galaxies will be measured, and then such calculations may provide a useful tool to distinguish between these two progenitor models (Perna \\& Belczynski 2001). Future and current space missions like HETE-II, INTEGRAL, GLAST or SWIFT will hopefully measure precise positions of a large number of bursts even of the short duration and settle down the issue of GRB progenitors." + }, + "0112/astro-ph0112408_arXiv.txt": { + "abstract": "We discuss the correlation between late-time integrated Sachs-Wolfe (ISW) effect in the cosmic microwave background (CMB) temperature anisotropies and the large scale structure of the local universe. This correlation has been proposed and studied in the literature as a probe of the dark energy and its physical properties. We consider a variety of large scale structure tracers suitable for a detection of the ISW effect via a cross-correlation. In addition to luminous sources, we suggest the use of tracers such as dark matter halos or galaxy clusters. A suitable catalog of mass selected halos for this purpose can be constructed with upcoming wide-field lensing and Sunyaev-Zel'dovich (SZ) effect surveys. With multifrequency data, the presence of the ISW-large scale structure correlation can also be investigated through a cross-correlation of the frequency cleaned SZ and CMB maps. While convergence maps constructed from lensing surveys of the large scale structure via galaxy ellipticities are less correlated with the ISW effect, lensing potentials that deflect CMB photons are strongly correlated and allow, probably, the best mechanism to study the ISW-large scale structure correlation with CMB data alone. ", + "introduction": "It is by now well known the importance of cosmic microwave background (CMB) temperature fluctuations as a probe of cosmology \\cite{est}. In addition to the dominant anisotropy contribution at the last scattering surface \\cite{Huetal97}, CMB photons, while on transit to us, encounter the large scale structure which imprints modifications on the temperature fluctuations. In general, large scale structure affects CMB through two distinct processes: gravity and Compton scattering. The modifications due to gravity arises from frequency changes via gravitational red and blue-shifts \\cite{SacWol67,ReeSci68}, through deflections involving lensing \\cite{Blaetal87,Hu00} and time-delays \\cite{HuCoo01}. During the reionized epoch, photos can both generate and erase primary fluctuations through scattering via free electrons \\cite{SunZel80,OstVis86}. Here, we discuss an effect due to gravitational redshift commonly known in the literature as the integrated Sachs-Wolfe (ISW; \\cite{SacWol67}) effect at late times. The temperature fluctuations in the ISW effect result from the differential redshift effect from photons climbing in and out of time evolving potential perturbations from last scattering surface to the present day. In currently popular cold dark matter cosmologies with a cosmological constant, significant contributions arise at redshifts of cosmological constant domination ($z \\lesssim 2$), at on and above the scale of the horizon at the time of decay. When projected on the sky, the ISW effect contributes at large angular scales and has a power spectrum that scales with the wavenumber as $k^{-5}$ times the linear density field power spectrum \\cite{HuSug94}. This is in contrast to most other contributions to CMB temperature fluctuations from the local universe, such as the well known thermal Sunyaev-Zel'dovich (SZ; \\cite{SunZel80}) effect, that peak at small angular scales and scales with the wavenumber as $k^{-1}$. Since time evolving potentials that contribute to the ISW effect may also be probed by observations of the large scale structure, it is then expected that the ISW effect may be correlated with certain tracers. The presence of the ISW effect can then be detected via a cross-correlation of the CMB temperature fluctuations at large angular scales and the fluctuations in an appropriate tracer field. Since the ISW contribution is sensitive to how one models cosmology at late-times, such as the presence of a dark energy component and its physical properties (for example, the ratio of dark energy pressure to density), the correlation between the CMB temperature and tracer fields has been widely discussed in the literature \\cite{CriTur96}. Though several attempts have already been made to cross-correlate the ISW effect, using the best COBE temperature map, and foreground sources, such as X-ray and radio galaxies, there is, so far, no clear detection of the correlation signal \\cite{Bouetal98}. As we discuss later, these non-detections are not surprising given the large sample variance associated with the correlating part of the temperature fluctuations, i.e., the ISW effect, with contribution to the variance coming from the primary temperature fluctuations. Even for the best case scenarios involving whole-sky observations and no noise contribution in the tracer field, the expected cumulative signal-to-noise for the ISW-large scale structure correlation is at most a ten. If the ISW-large scale structure correlation is to be used as probe of cosmological and astrophysical properties, it is certainly necessary to study in detail what tracers are best correlated with the ISW effect and why. In this paper, we address this question by studying in detail the correlation between the ISW effect and large scale structure. Since we are primarily interested in understanding what types of tracers are best suited to detect the correlation, we will consider a variety of large scale structure observations and tracers. These include luminous sources such as galaxies or active galactic nuclei (AGN) at different wavelengths, the dark matter halos or galaxy clusters that describe the large scale clustering of the universe, gravitational lensing, and other contributions to CMB from the large scale structure, such as the thermal SZ effect. Unlike previous studies \\cite{PeiSpe00,Veretal01}, we do not consider cosmological applications of the cross-correlation involving measurement of dark matter and dark energy properties. In \\S~\\ref{sec:method}, we outline the correlation between ISW effect and large scale structure. In \\S~\\ref{sec:discussion}, we discuss our results and study which tracer is best suited for correlation studies. We suggest that in addition to tracers involving sources such as galaxies or AGNs, dark matter halos that contain these sources may also be suitable for a detection of the ISW effect via a cross-correlation. The best probe of the ISW-large scale structure correlation is the lensing effect on CMB photons on transit to us from the last scattering surface. Since one can extract information related to lensing deflections from quadratic statistics on temperature data \\cite{SelZal99,Hu01}, this allows one to use all-sky CMB anisotropy maps, such as from the Planck surveyor, alone to extract the ISW-large scale structure correlation. ", + "conclusions": "\\begin{figure} \\centerline{\\psfig{file=iswgalsn.eps,width=3.6in,angle=-90}} \\caption{The signal-to-noise for the detection of source-ISW cross-correlation with a variety of sources that trace the large scale structure. The considered surveys include: HEAO X-ray catalog, radio sources from the NVSS, and the upcoming Sloan galaxy catalog. The signal-to-noise includes the fraction of sky covered in each of these surveys. The correlation with Sloan is limited by the sample variance associated with the ISW signal while the other two catalogs have significant shot-noise associated with the limited number counts.} \\label{fig:galsn} \\end{figure} \\begin{figure} \\centerline{\\psfig{file=iswhalosn.eps,width=3.6in,angle=-90}} \\caption{The signal-to-noise for the detection of source-ISW cross-correlation with dark matter halos that trace the large scale structure with low mass limits down to 10$^{15}$, 10$^{14}$ and 10$^{13}$ M$_{\\sun}$ (solid lines). The long-dashed line is the maximum signal-to-noise when the ISW effect can be perfectly separated from the dominant CMB data. The dot-dashed line shows an estimate for the correlation using a catalog of clusters based on the proposed Dark Universe Exploration Telescope (DUET).} \\label{fig:halosn} \\end{figure} We now discuss a variety of large scale structure tracers that can potentially be used to cross-correlate with CMB temperature data. \\subsection{ISW-source Correlation} Probably the most obvious tracer of the large scale structure density field in the linear regime is luminous sources such as galaxies at optical wavelengths and AGNs at X-rays and/or radio wavelengths. We can write the associated window function in this case as \\begin{equation} W^\\X(k,\\rad) = b_\\X(k,\\rad) n_\\X(\\rad) G(\\rad)\\, , \\end{equation} where $b_X(k,\\rad)$ is the scale-dependent source bias as a function of the radial distance and $n_X(\\rad)$ is the normalized redshift distribution of sources such that $\\int_0^{\\rad_0} n_\\X(\\rad)=1$. We will model the redshift distribution in current and upcoming source catalogs with an analytic form: \\begin{equation} n_\\X(z) = \\left(\\frac{z}{z_0}\\right)^{\\alpha} \\exp -\\left(\\frac{z}{z_0}\\right)^\\beta \\, , \\end{equation} where $(\\alpha,\\beta)$ denote the slope of the distribution at low and high $z$'s, respectively with a mean given by $\\sim z_0$. For the purpose of this calculation we take $\\alpha=\\beta=1$ and vary $z_0$ from 0.1 to 1.5 so as to mimic the expected sources from current and upcoming catalogs as well as to cover the redshift range in which ISW contributions are generally expected. Since any redshift dependence on bias can be included as a variation to the source redshift distribution, the scale dependence on bias is more important. Since the ISW effect is primarily associated with large linear scales, the source bias can be well approximated as scale independent. Such an assumption is fully consistent with results from numerical simulations \\cite{Benetal00}, results from redshift surveys \\cite{Veretal01} and semi-analytic calculations involving the so-called halo model \\cite{Sel00}. We take the source bias to be in the range of 1 to 3 as expected for galaxies and biased sources as X-ray objects. Since $C_l^{\\isw-\\X} \\propto b$, and $C_l^\\X \\propto b^2$, note that the correlation coefficient is independent of bias and other normalizing factors. The bias becomes only important for estimating the signal-to-noise. In order to estimate signal-to-noise, we describe the noise contribution associated with source catalogs by the finite number of sources one can effectively use to cross-correlate with temperature data. We can write this shot-noise contribution as \\begin{equation} C_l^{\\rm N_\\X} = \\frac{1}{\\bar{N}} \\end{equation} where $\\bar{N}$ is the surface density of source per steradian. In figure~\\ref{fig:galsn}, we show cumulative signal-to-noise for a cross-correlation of the temperature data with large scale structure. We have considered three source catalogs here involving the HEAO X-ray catalog ($f_sky=1/3$,$\\bar{N} \\sim 10^{3}$ sr$^{-1}$ and $b_\\X=3$), NVSS radio sources ($f_sky=0.82$,$\\bar{N}\\sim 2 \\times 10^{5}$ sr$^{-1}$ and $b_X=1.6$) \\cite{Conetal98} and galaxy counts down to R band magnitude of 25 from the Sloan Digital Sky survey ($f_\\sky=0.25$, $7 \\times 10^{8}$ sr$^{-1}$ and $b_\\X=1$) \\cite{Yoretal00}. The first two catalogs have already been used for this cross-correlation with the COBE data \\cite{Bouetal98}. The non-detection of the cross-correlation is not surprising given that we find cumulative signal-to-noise values less than 1; this estimate will become worse once the noise contribution associated with COBE temperature data is also included. In both cases, the shot-noise associated with tracer fields is significant. For Sloan, the limiting factor in the signal-to-noise is the large noise contribution associated with the ISW contribution due to dominant primary temperature fluctuations. Our estimates for the signal-to-noise for ISW-Sloan correlation is consistent with previous estimates \\cite{PeiSpe00}, where it was concluded that the detection may be challenging given the small signal-to-noise. In addition to luminous sources, one can also cross-correlate the temperature map with a catalog of galaxy clusters, or dark matter halos. It is expected that wide-field galaxy lensing surveys and small-angular resolution SZ surveys will allow mass selected catalogs of clusters \\cite{Holetal00}. Similar catalogs of clusters can also be compiled through wide-field galaxy catalogs such as Sloan and X-ray imaging data, such as from the proposed Dark Universe Exploration Telescope (DUET). The redshift distribution of halos in such catalogs follows simply from analytical arguments, such as through the mass function $dn(M,z)/dM$ calculated following analytical methods such as the Press-Schechter (PS; \\cite{PreSch74}) mass function or numerical measurements \\cite{Jenetal01}. Additionally, the halo bias is also well known through analytical methods \\cite{Moetal97}: \\begin{equation} b_h(M,z) = 1+ \\frac{\\left[\\nu^2(M,z) - 1\\right]}{\\delta_c}\\, , \\end{equation} where $\\nu(M,z) = \\delta_c/\\sigma(M,z)$ is the peak-height threshold. $\\sigma(M,z)$ is the rms fluctuation within a top-hat filter at the virial radius corresponding to mass $M$, and $\\delta_c$ is the threshold overdensity of spherical collapse. Useful fitting functions and additional information on these quantities could be found in \\cite{Hen00}. For the purpose of this calculation, we use $b_X = \\langle b_M \\rangle$, such that the mass averaged halo bias, as a function of redshift is: \\begin{equation} \\left< b_M \\right>(z) = \\frac{1}{\\bar{n}_h(z)} \\int_{M_{\\rm min}}^{\\infty} dM \\frac{dn(M,z)}{dM} b_h(M,z) \\, . \\end{equation} Here, the mean number density of halos, as a function of redshift, given by $\\bar{n}_h(z) = \\int dM dn(M,z)/dM$. The redshift distribution of halos follow similarly from PS theory: \\begin{equation} n_\\X(z) = \\frac{1}{\\bar{N}} \\frac{d^2V}{dzd\\Omega} \\left[ \\int_{M_{\\rm min}}^{\\infty} dM \\frac{dn(M,z)}{dM} \\right] \\, , \\end{equation} where \\begin{equation} \\bar{N} = \\int dz \\frac{d^2V}{dzd\\Omega} \\left[ \\int_{M_{\\rm min}}^{\\infty} dM \\frac{dn(M,z)}{dM} \\right] \\, , \\end{equation} with the comoving volume element given by $d^2V/dzd\\Omega$. Note that the shot-noise associated with the halo-halo power spectrum is given by the surface-density of halos $C_l^{\\rm N_\\X} \\equiv 1/\\bar{N}$. Note that even though the surface density of halos are significantly smaller than, say, the case for galaxies or sources, one gains an equivalent factor with the increase in halo bias relative to the same for galaxies. Therefore, on average, we expect catalogs based on halos to produce same order of magnitude signal-to-noise as for luminous sources. In figure~\\ref{fig:halosn}, we show the cumulative signal-to-noise for detecting the ISW-halo correlation. Here, we assume all-sky catalogs of clusters down to a minimum mass limit of 10$^{13}$, 10$^{14}$ and 10$^{15}$ M$_{\\sun}$. The lower limit of 10$^{14}$ M$_{\\sun}$ is consistent with the expected mass threshold for SZ clusters that will be detected with the planned South Pole Telescope \\cite{Holetal00}. This mass limit is constant over a wide range in redshift, while mass catalogs based on lensing surveys contain a redshift dependent minimum mass limit due to the variation in the lensing window function. The decrease in signal-to-noise with increase in mass is due to the decrease in the surface density of halos and thus the increase in the shot-noise associated with the halo side of the correlation. Though the halo surface density decreases as a function of mass, this is partly accounted by the increase in the halo bias such that even with a low surface density of halos, one can attempt a correlation of the temperature with cluster catalogs. As an example of a temperature-source catalog correlation that can be expected with the MAP data, we also show an estimate for the signal-to-noise for a catalog of $\\sim$ 18000 clusters that is expected to be compiled from the wide-field X-ray survey by the proposed DUET mission (dot-dashed line). Though the cumulative signal-to-noise is below 2, the DUET catalog is significantly preferred over the Sloan survey due to the fact that cluster bias associated with DUET tracers can be a priori known through mass estimates based on the electron temperature data while galaxy bias may be more complicated. \\begin{figure*} \\centerline{\\psfig{file=iswsz.eps,width=3.6in,angle=-90} \\psfig{file=szisw.eps,width=3.6in,angle=-90}} \\caption{{\\it Left:} The cross-correlation power spectrum between the ISW and SZ effects. For comparison, we also show the power spectra of ISW and SZ effects and the CMB anisotropies. {\\it Right:} The cumulative signal-to-noise for the detection of the ISW-SZ correlation with Planck and MAP data and using spectral dependence of the SZ contribution to separate it out from thermal CMB fluctuations. The maximum signal-to-noise is when the ISW effect is separated from the dominant primary fluctuations at the last scattering surface.} \\label{fig:iswsz} \\end{figure*} \\subsection{ISW-SZ Correlation} Following the derivation of the ISW-large scale structure correlation, we can also consider the cross-correlation between the ISW effect and the Sunyave-Zel'dovich thermal effect \\cite{SunZel80} due to inverse-Compton scattering of CMB photons via hot electrons. The temperature decrement due to the SZ effect can be written as the integral of pressure along the line of sight \\begin{equation} y\\equiv\\frac{\\Delta T}{T_{\\rm CMB}} = g(x) \\int d\\rad a(\\rad) \\frac{k_B \\sigma_T}{m_e c^2} n_e(\\rad) T_e(\\rad) \\, \\end{equation} where $\\sigma_T$ is the Thomson cross-section, $n_e$ is the electron number density, $\\rad$ is the comoving distance, and $g(x)=x{\\rm coth}(x/2) -4$ with $x=h \\nu/k_B T_{\\rm CMB}$ is the spectral shape of SZ effect. At Rayleigh-Jeans (RJ) part of the CMB, $g(x)=-2$. For the rest of this paper, we assume observations in the Rayleigh-Jeans regime of the spectrum. Due to the spectral dependence of the SZ effect when compared to CMB thermal fluctuations, the SZ signal can be extracted from CMB fluctuations in multifrequency data \\cite{Cooetal00}. Here, we use expected results from such a frequency separation with Planck and MAP data, and consider the ISW-SZ cross-correlation by correlating the CMB and SZ maps. When calculating signal-to-noise, we will use the noise power spectra calculated in \\cite{Cooetal00} for the Planck SZ and CMB maps. Following our discussion on the ISW-large scale structure correlation, we can write the relevant window function for the SZ effect as \\begin{equation} W^{\\rm SZ}(r) = g(x) \\frac{k_B \\sigma_T b_g(\\rad) \\bar{n}_e}{a(r)^2 m_e c^2} \\label{eqn:wsz} \\end{equation} where, at linear scales corresponding to the ISW effect, the pressure bias, $b_g$, relative the density field follows from arguments based on the halo approach to large scale pressure fluctuations \\cite{Sel00,Coo01}: \\begin{equation} b_g(z) = \\int dM \\frac{M}{\\bar{\\rho}}\\frac{dn(M,z)}{dM} b_{\\rm halo}(M,z) T_e(M,z) \\, . \\end{equation} Here, $T_e(M,z)$ is the electron temperature, which can be calculated through arguments related to the virial theorem. In Eq.~\\ref{eqn:wsz}, $\\bar{n}_e$ is the mean density of electrons today. In figure~\\ref{fig:iswsz}, we show the cross-correlation between the ISW and SZ effects. For comparison, we also show the ISW and SZ power spectra. The correlation coefficient for the ISW-SZ effect ranges from 0.3, at $l \\sim$ few tens to 0.1 at $l \\sim$ few hundred suggesting that ISW and the SZ contributions are not strongly correlated. This could be understood based on the fact that contributions to the SZ effect primarily comes from the so-called 1-halo term of the halo models of non-linear clustering and not the 2-halo term that tracers the large scale correlations and, thus, the linear density fluctuations responsible for the ISW effect. In the same figure, we also show the signal-to-noise for the detection of the ISW-SZ cross-correlation. For an experiment like Planck, we find that the signal-to-noise ratio is at the level of $\\sim$ 4, while for MAP, it is at the level of 0.3; this is understandable as MAP has no high frequency information for a reliable separation of the SZ effect. In order to investigate what limits the signal-to-noise for the detection of the ISW-SZ correlation with a mission like Planck, we decided to set the noise contribution to the ISW effect as simply due to the ISW effect itself, instead of the total CMB power spectrum. This led to the dot-dashed line. Further removing the noise contribution to the SZ map, such that a perfect separation of the SZ effect is possible led to the dotted line, which converges to the previous dot-dashed line by $l \\sim 10$. It is clear that the dominant noise contribution comes from the noise associated with the ISW effect, and effectively, the dominant CMB anisotropies at the last scattering which cannot be easily separated from the ISW effect. If there is a useful separation scheme to extract the ISW contribution alone, the expected signal-to-noise for the ISW-SZ correlation is at the level of $\\sim$ 60 suggesting a clear detection of the correlation which can be used for cosmological and astrophysical purposes. \\begin{figure} \\centerline{\\psfig{file=convergence.eps,width=3.6in,angle=-90}} \\caption{The power spectrum of convergence constructed from CMB deflections (top curve) and galaxy shape data (bottom curve). In the case of reconstruction based on CMB, we show expected errors from the Planck mission while for large scale structure weak lensing, we show expected errors for a survey of 400 deg.$^2$ down to a R band magnitude of 25.} \\label{fig:lensing} \\end{figure} \\subsection{ISW-lensing correlation} We can consider two forms of lensing: the potentials that deflect CMB photons and potentials that shear background galaxy images. The former can be constructed using quadratic statistics in temperature data or using Fourier space statistics that are optimized to extract the lensing signal, while the latter is probed in weak lensing survey using galaxy shapes. The deflection angle of CMB photons on the sky, $\\alpha(\\bn) = \\nabla \\phi(\\bn)$, are given by the gradient of the projected potential $\\Phi$ (see e.g. \\cite{Kai92}), \\begin{eqnarray} \\phi(\\bm)&=&- 2 \\int_0^{\\rad_0} d\\rad \\frac{\\da(\\rad_0-\\rad)}{\\da(\\rad)\\da(\\rad_0)} \\Phi (\\rad,\\hat{{\\bf m}}\\rad ) \\,. \\label{eqn:lenspotential} \\end{eqnarray} The lensing potential can be related to the well known convergence generally encountered in conventional lensing studies involving galaxy shear \\begin{eqnarray} \\kappa(\\bm) & =& {1 \\over 2} \\nabla^2 \\phi(\\bm) \\nonumber \\\\ & = &-\\int_0^{\\rad_0} d\\rad \\frac{\\da(\\rad)\\da(\\rad_0-\\rad)}{\\da(\\rad_0)} \\nabla_{\\perp}^2 \\Phi (\\rad ,\\hat{{\\bf m}}\\rad) \\, , \\nonumber \\\\ \\end{eqnarray} where note that the 2D Laplacian operating on $\\Phi$ is a spatial and not an angular Laplacian. Expanding the lensing potential to Fourier moments, \\begin{equation} \\phi(\\bn) = \\int \\frac{d^2\\vecl}{(2\\pi)^2} \\phi(\\vecl) {\\rm e}^{i \\vecl \\cdot \\bn} \\, , \\end{equation} we can write the usually familiar quantities of convergence and shear components of weak lensing as \\cite{Hu00} \\begin{eqnarray} \\kappa(\\bn) &=& -\\frac{1}{2}\\int \\frac{d^2\\vecl}{(2\\pi)^2} l^2 \\phi(\\vecl) {\\rm e}^{i\\vecl \\cdot \\bn} \\nonumber \\\\ \\gamma_1(\\bn) \\pm i\\gamma_2(\\bn) &=& -\\frac{1}{2}\\int \\frac{d^2\\vecl}{(2\\pi)^2}\\ l^2 \\phi(\\vecl) {\\rm e}^{\\pm i 2 (\\phi_l-\\phi)}{\\rm e}^{i\\vecl \\cdot \\bn} \\, . \\label{eqn:kappaphi} \\end{eqnarray} Though the two terms $\\kappa$ and $\\phi$ contain differences with respect to radial and wavenumber weights, these differences cancel with the Limber approximation \\cite{Lim54}. In particular, their spherical harmonic moments are simply proportional \\begin{eqnarray} \\phi_{l m} &=& -{2 \\over l(l+1)} \\kappa_{l m} = \\int d {\\bn} \\Ylmn{}^*(\\bn) \\phi(\\bn) \\nonumber\\\\ &=& i^l \\int {d^3 {\\bf k}\\over 2\\pi^2} \\delta({\\bf k}) \\Ylmn{}^* (\\bk) I_\\ell^{\\rm len}(k) \\label{eqn:GSSZequiv} \\end{eqnarray} with \\begin{eqnarray} I_\\ell^{\\rm len}(k)& =& \\int W^{\\rm len}(k,r) j_l(k\\rad) \\,,\\nonumber\\\\ W^{\\rm len}(k,r)& =& -3 \\frac{\\Omega_m}{a} \\left({H_0 \\over k}\\right)^2 G(r) {\\da(\\rad_0 - \\rad) \\over \\da(\\rad)\\da(\\rad_0)}\\,. \\label{eqn:lensint} \\end{eqnarray} Here, we have used the Rayleigh expansion of a plane wave, equation~(\\ref{eqn:Rayleigh}), and the fact that $\\nabla^2 \\Ylmn = -l(l+1) \\Ylmn$. \\begin{figure} \\centerline{\\psfig{file=kappasn.eps,width=3.6in,angle=-90}} \\caption{The signal-to-noise for the detection of lensing-ISW cross-correlation with galaxy shear data and using deflections in the CMB. The dotted lines are for galaxy weak lensing surveys with background sources at redshifts of 0.5 and 1.5, respectively, and with a sky area of $\\pi$ steradians, similar to the Sloan survey. The dashed lines show the signal-to-noise when temperature data is cross-correlated with an estimator for lensing deflections in Planck and MAP temperature data. The solid line is when a temperature map is cross-correlated with a noise-free estimator of deflections.} \\label{fig:lensingsn} \\end{figure} For the construction of deflection angles based on the CMB temperature data, we make use the quadratic statistic proposed by \\cite{Hu01} involving the divergence of the temperature weighted temperature gradients, $\\nabla \\cdot (\\theta \\nabla \\theta)$. In Fourier space, we can write the estimator for the deflection angles as \\begin{eqnarray} D(\\vecl) &=& \\frac{N_l}{l} \\int \\frac{d^2\\vecl_1}{(2\\pi)^2} \\left(\\vecl \\cdot \\vecl_1 C_{l_1}^{\\rm CMB} +\\vecl \\cdot (\\vecl-\\vecl_1) C_{|\\vecl-\\vecl_1|}^{\\rm CMB}\\right)\\nonumber \\\\ && \\quad \\times \\frac{\\theta(l_1)\\theta(|\\vecl-\\vecl_1|)}{2 C_{l_1}^{\\rm tot} C_{|\\vecl-\\vecl_1|}^{\\rm tot}} \\, . \\end{eqnarray} The ensemble average, $\\langle D(\\vecl) \\rangle$, is equal to the deflection angle, $l\\phi(\\vecl)$, when \\begin{equation} N_l^{-1} = \\frac{1}{l^2} \\int \\frac{d^2\\vecl_1}{(2\\pi)^2} \\frac{\\left(\\vecl \\cdot \\vecl_1 C_{l_1}^{\\rm CMB} +\\vecl \\cdot (\\vecl-\\vecl_1) C_{|\\vecl-\\vecl_1|}^{\\rm CMB}\\right)^2}{2 C_{l_1}^{\\rm tot} C_{|\\vecl-\\vecl_1|}^{\\rm tot}} \\, . \\end{equation} Note that $N_l$ is the noise power spectrum associated with the reconstructed deflection angle power spectrum: \\begin{equation} \\langle D(\\vecl) D(\\vecl') \\rangle = (2\\pi)^2 \\delta_D(\\vecl+\\vecl') \\left(l^2 C_l^{\\phi \\phi} +N_l\\right) \\, . \\end{equation} In the case of lensing surveys using galaxy shear data, we rewrite equations~(\\ref{eqn:kappaphi}) and (\\ref{eqn:lensint}) such that $\\da(\\rad_0) = \\da(\\rad_s)$ where $\\rad_s)$ is radial distance to background sources from which shape measurements are made. We assume that all sources are at the same redshift, though,a distribution of sources in the redshift range expected does not lead to a significantly different result than the one suggested here. The shot-noise contribution to the convergence power spectrum associated with lensing surveys involving galaxy ellipticity data is \\begin{equation} C_l^{\\rm N_\\X} = \\frac{\\langle \\gamma_{\\rm int}^2 \\rangle}{\\bar{n}} \\, , \\end{equation} where $\\langle \\gamma_{\\rm int}^2 \\rangle^{1/2}$ is the rms noise per component introduced by intrinsic ellipticities, typically $\\sim$ 0.6 for best ground based surveys, and $\\bar{n}$ is the surface number density of background source galaxies from which shape measurements can be made. For surveys that reach a limiting magnitude in $R \\sim 25$, the surface density is consistent with $\\bar{n} \\sim 6.9 \\times 10^8$ sr$^{-1}$ or $\\approx 56$ gal arcmin$^{-2}$ \\cite{Smaetal95}, such that $C_l^{\\rm N} \\sim 2.3 \\times 10^{-10}$. In figure~\\ref{fig:lensing}, we compare the lensing convergence power spectrum associated with CMB (top curve) and a large scale structure weak lensing survey from galaxy ellipticities with background source galaxies at a redshift of one. Note that we have obtained the convergence power spectrum associated with lensing deflections in CMB following the estimator for the lensing deflection power spectrum and the two are simply related following equation~(\\ref{eqn:kappaphi}) such that $C_l^{\\kappa} = l^4/4 C_l^{\\phi\\phi}$. For comparison, we also show expected error bars on the reconstructed convergence power spectrum from CMB, via the Planck temperature data, and for a wide-field survey of 400 deg.$^2$ down to a R-band magnitude of 25. In figure~\\ref{fig:lensingsn}, we show the associated cumulative signal-to-noise in the detection of the ISW-lensing correlation for galaxy lensing surveys and using CMB data. We assume an area of $\\pi$ steradians for the lensing surveys and as the signal-to-noise scales as $f_{\\rm sky}^{1/2}$, we do not expect a significant use of the current and upcoming lensing surveys which are restricted to at most few hundred sqr. degrees. The dedicated instruments, such as the Large-aperture Synoptic Survey Telescope (LSST; \\cite{TysAng00}) however, will provide wide-area maps of the lensing convergence and these will certainly be useful for cross-correlation studies with CMB to extract the ISW effect. Note that the Planck data allow the best opportunity to detect the ISW effect by correlating an estimator for deflections with a temperature map. The MAP has a lower cumulative signal-to-noise due to its estimator of deflections is affected by the low resolution of the temperature data. Nevertheless, the MAP data will certainly allow the first opportunity to detect the presence of the ISW effect either from CMB data alone or through cross-correlation of other tracers. \\begin{figure} \\centerline{\\psfig{file=corr.eps,width=3.6in,angle=-90}} \\caption{The correlation-coefficient for ISW-large scale structure correlations involving lensing effect on CMB (solid line), a catalog of dark matter halos down to a mass limit of $10^{14}$ M$_{\\sun}$ at all redshifts (dotted line), lensing convergence from galaxy shear data (dashed lines) with sources at redshifts of 0.5 and 1.5, and sources as tracers with mean redshifts of 0.4, 0.7 and 1.3. The potentials that deflect CMB photons and sources at $z \\sim 1.5$ are best correlated with the ISW effect.} \\label{fig:corr} \\end{figure} \\begin{figure} \\centerline{\\psfig{file=iswkspace.eps,width=3.6in,angle=-90}} \\caption{The contribution to the power spectra of ISW effect and other tracers as indicated. The sources have a mean redshift, $z_0$ of 0.4, 0.7 and 1.3 (from top to bottom). The plot shows that the contribution to the lensing effect on CMB comes from same Fourier modes as the ISW effect while there is a mismatch when compared to galaxy surveys at low redshifts.} \\label{fig:kspace} \\end{figure}" + }, + "0112/astro-ph0112314_arXiv.txt": { + "abstract": "The recent developments in studies of TeV radiation from {\\em blazars} are highlighted and the implications of these results for derivation of cosmologically important information about the {\\em cosmic infrared background radiation} are discussed. ", + "introduction": "Since early 90's many papers have been published with an unusual combination of two keywords - {\\em Blazars} and {\\em Cosmic-Infrared-Background} (CIB) radiation. Formally, one may argue that there is no apparent link between these two topics. \\noindent $\\bullet$ The blazars constitute a sub-class of AGN dominated by highly variable (several hours or less) components of broadband (from radio to gamma-rays) {\\sl non-thermal emission} produced in relativistic jets pointing close to the line of sight. \\noindent $\\bullet$ CIB is a part of the overall diffuse extragalactic background radiation (DEBRA) dominated by {\\sl thermal emission} components produced by stars and dust, and accumulated over the entire history of the Universe. \\vspace{1mm} While the blazars may serve as ideal laboratories for study of MHD structures and particle acceleration processes in relativistic jets, CIB carries crucial cosmological information about the formation epochs and history of evolution of galaxies. To a large extent, these two topics are relevant to quite independent areas of modern astrophysics and cosmology. Yet, the current studies of CIB and blazars, more specifically the sub-population of blazars emitting TeV gamma-rays ({\\em TeV blazars}), are tightly coupled through the intergalactic (IG) absorption of TeV radiation by infrared photons of DEBRA. The astrophysical/cosmological importance of this interesting effect (Nikishov, 1962; Gould and Schreder, 1966; Jelly, 1966; Stecker et al., 1992) was clearly recognized after the discovery of TeV $\\gamma$-rays from two BL Lac objects -- Mkn~421 and Mkn~501 (for review see e.g. Vassiliev 2000). ", + "conclusions": "The energy-dependent mean free path of $\\gamma$-rays in the intergalactic medium at TeV energies does not exceed several 100 Mpc. Therefore VHE $\\gamma$-rays from TeV blazars arrive with significantly distorted spectra. Our limited knowledge of CIB results in large uncertainties in the reconstructed (corrected for intergalactic absorption) intrinsic $\\gamma$-ray spectra. Consequently, in spite of good quality of spectral measurements of two strongest TeV blazars, Mkn 421 and Mkn 501, nowadays we are faced with a challenge - our understanding of radiation processes in relativistic jets is neither complete nor conclusive. Nevertheless, the gamma-ray astronomers believe that eventually (soon ?) they will learn to identify confidently the radiation mechanisms, to fix/constrain the relevant model parameter space, and calculate robustly the intrinsic $\\gamma$-ray spectra based on the multiwavelength studies of spectral and temporal characteristics of blazars obtained simultaneously at X-ray and TeV bands on sub-hour timescales. Then it will be possible to estimate unambiguously the effect of the intergalactic $\\gamma$-ray absorption, and thus to infer robust information about the CIB fluxes. Moreover, the studies of angular and spectral properties of giant Pair Halos formed around powerful extragalactic multi-TeV sources may provide us with a unique tool for derivation of the spectra and absolute fluxes of CIB at {\\sl different cosmological epochs z}, and thus to probe the evolution of galaxies in past." + }, + "0112/astro-ph0112064_arXiv.txt": { + "abstract": "We present the results of a joint investigation aimed at constraining the primordial He content ($Y_P$) on the basis of both the Cosmic Microwave Background (CMB) anisotropy and two stellar observables, namely the tip of the Red Giant Branch (TRGB) and the luminosity of the Zero Age Horizontal Branch (ZAHB). Current baryon density estimates based on CMB measurements cover a wide range values $0.009\\la \\Omega_bh^2 \\la 0.045$, that according to Big Bang Nucleosynthesis (BBN) models would imply $0.24\\la Y_P \\la 0.26$. We constructed several sets of evolutionary tracks and HB models by adopting $Y_P=0.26$ and several metal contents. The comparison between theory and observations suggests that ZAHB magnitudes based on He-enhanced models are 1.5$\\sigma$ brighter than the empirical ones. The same outcome applies for the TRGB bolometric magnitudes. This finding somewhat supports a $Y_P$ abundance close to the canonical 0.23-0.24 value. More quantitative constraints on this parameter are hampered by the fact that the CMB pattern shows a sizable dependence on both $Y_P$ and the baryon density only at small angular scales, i.e. at high $l$ in the power spectrum ($l\\ga 100$). However, this region of the power spectrum could be still affected by deceptive systematic uncertainties. Finally, we suggest to use the {\\em UV-upturn} to estimate the He content on Gpc scales. In fact, we find that a strong increase in $Y_P$ causes in metal-poor, hot HB structures a decrease in the UV emission. ", + "introduction": "The comparison between chemical abundances of deuterium, helium, and lithium predicted by BBN models with current empirical estimates is one of the most viable method to constrain the physical mechanisms and the cosmology which governed the nucleosynthesis of primordial abundances (Olive, Steigman, \\& Walker 2000). As far as the primordial He content is concerned, current empirical estimates are mainly based on measurements of nebular emission lines in low-metallicity, extragalactic HII regions (Izotov, Thuan, \\& Lipovetsky 1997; Olive, Steigman, \\& Skillman 1997). Recent He determinations present small observational errors ($\\approx 1$\\%), but large uncertainties between independent measurements: $Y_P=0.234\\pm0.003$ by Olive \\& Steigman (1995) against $Y_P=0.244\\pm0.002$ by Izotov \\& Thuan (1998). This evidence suggests that current He abundances are still dominated by systematic errors. In fact, Viegas, Gruenwald, \\& Steigman (2000) and Gruenwald, Steigman, \\& Viegas (2001) in two detailed investigations on the ionization correction for unseen neutral and doubly-ionized He in HII regions, found that He estimates should be reduced by 0.006 ($Y_P=0.238\\pm0.003$), a quantity which is a factor of 2-3 larger than typical statistical errors quoted in the literature. Moreover and even more importantly, Pistinner et al. (1999) on the basis of a new grid of stellar atmosphere models for OB stars found that the inclusion of both NLTE and metal-line blanketing effects causes an increase of the order of 40\\% in the ratio of He to H ionizing photons. This evidence together with uncertainties due to the occurrence of stellar winds, shocks, temperature fluctuations (Izotov, Thuan, \\& Lipovetsky 1997; Pistinner et al. 1999; Peimbert, Peimbert, \\& Luridiana 2001; Sauer, \\& Jedamzik 2001, and references therein) and of peculiar nebular dynamics certainly affects the He abundance estimates based on giant extragalactic HII regions. In addition it is worth mentioning that the HII regions used for determining the cosmological Helium abundance could have been somewhat polluted by the stellar yields of the pristine type II Supernovae, and in turn the empirical He abundances in these stellar systems should be corrected for self-pollution by massive stars. A plain evidence of this occurrence has been recently provided by Aloisi, Tosi, \\& Greggio (1999), and \\\"Ostlin (2000). On the basis of deep HST optical and NICMOS data they have resolved the stellar content of I~ZW~18 and found evidence that this blue compact galaxy hosts a relatively old population of asymptotic giant branch stars ($\\approx$ 0.1-5 Gyr). On the other hand, the comparison between star counts of horizontal branch (HB, central He burning phase) and red giant (RG, H shell burning phase) stars in Galactic Globular Clusters (GGCs) with the lifetimes predicted by evolutionary models, the so-called R parameter (Iben 1968), supplies upper limits to primordial He mass fraction of the order of 0.20 (Sandquist 2000; Zoccali et al. 2000). However, such estimates should be cautiously treated (Bono et al. 1995; Cassisi et al. 1998), since they are hampered by current uncertainties on the nuclear cross-section of the $^{12}C(\\alpha,\\gamma)^{16}O$ reaction (Buchmann 1996). Note that spectroscopic measurements of He abundances in low-mass population II stars are useless for constraining the primordial He content, because the He lines are either too faint (low-temperature stars) or affected by gravitational settling such as high temperature HB stars (Giannone \\& Rossi 1981; Moheler et al. 1999). However, empirical and statistical errors affecting abundance determinations of primordial deuterium, $^3$He, and lithium could be significantly larger than for He (Sasselov \\& Goldwirth 1995; Olive et al. 2000). Moreover, the primordial He content plays a paramount role in constraining both stellar ages and cosmic distances, since the Mass-Luminosity (M/L) relation of low and intermediate-mass stars during H and He burning phases depends on $Y_P$ (Bono et al. 2000). At the same time, at fixed He to metal enrichment ratio the He abundance adopted to model evolutionary and pulsational properties of metal-rich stellar structures does depend on $Y_P$ as well (Bono et al. 1997; Zoccali et al. 2000). The physical baryon density of the universe is one of the observables that can be determined with high accuracy using measurements of CMB anisotropies at intermediate and small angular scales (see e.g., Hu et al. 2000, and references therein). It goes without saying that this observable plays a key role not only to assess the plausibility of the physical assumptions adopted in BBN models (Tegmark \\& Zaldarriaga 2000) but also for constraining the intrinsic accuracy of current primordial abundance estimates. According to the joint analysis of both BOOMERanG and MAXIMA-1 data, it has been estimated at 68\\% confidence level a baryon density $\\Omega_B h^2= 0.032_{-0.004}^{+0.005}$ (Jaffe et al. 2001). On the basis of this observable Esposito et al. (2001) found that the new CMB measurements are inconsistent at more that 3$\\sigma$ with both standard and degenerate BBN models. On the other hand, the latest analysis of the BOOMERanG results, which has improved the removal of systematics from the data (Netterfield et al. 2001), found $\\Omega_b h^2=0.022^{+0.004}_{-0.003}$ (de Bernardis et al. 2001), in very good agreement with the BBN value. The same conclusion has been derived from the analysis of the ground-based CMB observations performed by the DASI interferometer (Halverson et al. 2001), which also found $\\Omega_b h^2=0.022^{+0.004}_{-0.003}$ (Pryke et al. 2001). This notwithstanding, the new analysis of the MAXIMA data (Lee et al. 2001), which extended the high $l$ coverage of the power spectrum measurement, still points towards somewhat higher values of the physical baryon density: $\\Omega_b h^2=0.032\\pm 0.006$ (Stompor et al. 2001). Finally, we mention the measurements of the CMB power spectrum at $l > $1000 by the Cosmic Background Imager. From these observations, and by assuming a flat cosmological model, the likelihood for the physical baryon density is found to peak at $\\Omega_b h^2$=0.009 (Padin et al. 2001). Current physical baryon densities based on CMB measurements and BBN models would imply that $Y_P$ might range from roughly 0.24 to approximately 0.26. The main aim of this investigation is to constrain $Y_P$ on the basis of recent CMB measurements and two stellar observables that depend on $Y_P$, namely the TRGB luminosity and the ZAHB luminosity. In \\S 2 we discuss in detail the adopted theoretical framework as well as the comparison between predicted and empirical observables. The effect of a change in $Y_P$ abundance on the {\\em UV-upturn} as well as on CMB anisotropies are presented in \\S 3 and \\S 3.1 respectively. Our conclusions and final remarks are briefly mentioned in \\S 4. ", + "conclusions": "Recent measurements of the CMB anisotropy provided the unique opportunity to evaluate several fundamental cosmological parameters and to supply for all of them a preliminary but plausible estimate of their error budget. The impact of these new measurements on cosmological models uncorked a flourishing literature. However, it is not easy to assess on a quantitative basis to what extent current differences in the physical baryon density derived from CMB observations are caused by deceptive systematic errors. As a matter of fact, CMB measurements of $\\Omega_bh^2$ range from 0.009 (Padin et al. 2001) to $0.032\\pm0.012$ (95\\% confidence level, Stompor et al. 2001). According to BBN models the new measurements imply that $Y_P$ might range from approximately 0.24 to roughly 0.26. By adopting the upper limit on $Y_P$ we investigated the impact of the change on two stellar observables, namely the ZAHB luminosity and the luminosity of the tip of the RGB. The main outcome of our analysis is that an increase in the primordial He content from the canonical $Y_P=0.23$ to $Y_P=0.26$ does not seem to be supported by the comparison between current theoretical predictions and empirical data. We found that the {\\em UV-upturn} can be adopted to estimate the primordial He content. In fact, numerical experiments suggest that an increase of $Y_P$ from 0.23 to 0.50 causes a decrease in the UV emission at least of the order of 20\\%. This is a preliminary rough estimate based on the assumption that AGB-manqu\\`e structures are the main sources of the {\\em UV-upturn}. An interesting feature of this observable is that current instruments can allow us to measure the {\\em UV-upturn} up to distances of the order of Gpcs. Note that to supply quantitative estimates of $Y_P$ on the basis of the comparison between synthetic and observed {\\em UV-upturns} it is necessary to account for the SED typical of complex stellar populations as a function of redshift (Tantalo et al. 1996; Yi et al. 1999). However, theoretical predictions should be cautiously treated, since UV flux when moving from low to high metal contents strongly depends on the efficiency of the mass loss as well as on the He to metal enrichment ratio (Greggio \\& Renzini 1999). The scenario has been further complicated by recent spectroscopic measurements of hot HB stars ($T_e \\ge 10,000$ K) in metal-poor GGCs (M15, M13). In fact, Behr et al. (2000) and Behr, Cohen, \\& McCarthy (2000) found that in these stars the iron abundance is enhanced by 1-2 order of magnitudes, whereas the He content is depleted by at least one order of magnitude respect to solar abundance. Unfortunately, we still lack quantitative estimates of the impact that such a peculiarities have on the UV emission. Theoretical and empirical arguments support the evidence that the density of baryons in the universe is homogeneous (Copi, Olive, \\& Schramm 1995). The same outcome applies to large scale chemical inhomogeneities (Copi, Olive, \\& Schramm 1996). However, it has been recently suggested by Dolgov \\& Pagel (1999, hereinafter DP) a new cosmological model that predicts a substantial spatial variation in the primordial chemical composition and a small baryon density variation. This investigation was triggered by a difference of one order of magnitude in the deuterium abundance of damped $Ly\\alpha$ systems along the line of sight of high-redshift ($0.5 \\le z \\le 3.5$) QSOs (D'Odorico et al. 2001; Steigman et al. 2001). The scenario developed by DP relies on a model of leptogenesis (Dolgov 1992) in which takes place a large lepton asymmetry and this asymmetry undergoes strong changes on spatial scales ranging from Mega to Giga pcs. The key feature of this model is to predict a large and varying lepton asymmetry and a small baryon asymmetry. Within this theoretical framework the He mass fraction in deuterium-rich regions should range from 35\\% to 60\\%, while the $Li$ one should increase up to $10^{-9}$, while the variation of the photon temperature should be $\\delta T/T\\approx 2.5\\times10^{-3}$. Obviously, the hypothesis that current changes in baryon density are due to real spatial variations is premature as any further speculative issue. Future full-sky CMB observations from space missions such as NASA's MAP (Wright 1999) and ESA's PLANCK (Mandolesi et al. 1998) will play a crucial role to properly address the problem of the spatial variation, since they will supply a larger sensitivity up to very high $l$ ($l > 1000$) and an improved control on systematics." + }, + "0112/astro-ph0112252_arXiv.txt": { + "abstract": "We summarize some of the major results obtained so far from the REFLEX survey of X-ray clusters of galaxies, concentrating on the latest measurements of the cluster X-ray luminosity function and two-point correlation function. The REFLEX luminosity function provides the most homogeneous census of the distribution function of masses in the local Universe, representing a unique zero-redshift reference quantity for evolutionary studies. On the other hand, the observed clustering of REFLEX clusters is very well described by the correlation function of a low-$\\Omega_M$ CDM model. Also, the bidimensional correlation map $\\xi(r_p, \\pi)$ shows no stretching along the line of sight, indicating negligible spurious effects in the sample, with at the same time a clear compression of the contours as expected in the presence of coherent motions. ", + "introduction": "The REFLEX\\footnote{The REFLEX Team includes: H. B\\\"ohringer (MPE), L. Guzzo (OAB), C.A. Collins (LJMU), P. Schuecker (MPE), G. Chincarini (OAB), R. Cruddace (NRAL), A.C. Edge (Durham), S. De Grandi (OAB), D.M. Neumann (Saclay), T. Reiprich (MPE), S. Schindler (LJMU), P.A. Shaver (ESO), W. Voges (MPE)} (ROSAT-ESO Flux-Limited X-ray) cluster survey is currently the largest sample of X-ray selected clusters of galaxies from the ROSAT All-Sky Survey with 1) a statistically homogeneous and fairly well-understood selection function; 2) measured redshifts. The survey contains 452 clusters over the southern celestial hemisphere ($\\delta<2.5^\\circ$), at galactic latitudes $|b_{II}|>20^\\circ$ and is more than 90\\% complete to a flux limit of $3 \\times 10^{-12}$ erg s$^{-1}$ cm$^{-2}$ (in the ROSAT band, 0.1--2.4 keV). X-ray fluxes and source extensions were re-measured from the RASS with a dedicated algorithm, thus avoiding the limitations of the standard analysis software in the characterisation of extended sources. The details of the whole identification process and a critical discussion of potential sources of incompleteness are presented in B\\\"ohringer et al. (2001a). Redshifts for all but 3 REFLEX clusters have been measured during a long Key Programme (1992-2000) using ESO telescopes (e.g. Guzzo et al. 1999), that collected more than 3500 galaxy redshifts over almost 500 X-ray targets. \\begin{figure} \\plotfiddle{guzzofig1.ps}{6.0cm}{0}{40}{40}{-150}{-40} \\caption{The sky distribution of REFLEX clusters. The dotted band marks the region of the galactic plane ($|b_{II}|>20^\\circ$) which is excluded from the survey.} \\label{aitoff} \\end{figure} Fig.~1 shows the distribution on the sky of the 452 clusters in the REFLEX survey, centered on the South Galactic Cap region (the largest contiguous area covered), while Fig.~2 plots their X-ray luminosity against their redshift. Given the survey flux limit (defined by the lower boundary of the point distribution), at large redshifts we are allowed to detect only the very bright, massive clusters. On the other hand, the very large solid angle of the survey allows for an extremely large volume to be explored (4.24 steradians, corresponding to $8.7\\times 10^8$ h$^{-3}$ Mpc$^3$ out to $z=0.3$ in an $\\Omega_M=0.3$, $\\Omega_\\Lambda=0.7$ cosmology), such that these extremely rare objects on the tail of the cluster X-ray luminosity function have a significant probability to be detected. It is not by chance, in fact, that the most luminous cluster known is still that discovered by the REFLEX survey in 1994, i.e. RXCJ1347.4-1144 (Schindler et al. 1995). Another example is shown in Fig.~4. Such a large volume is also ideal for sampling the very large modes of density fluctuations in the Universe (Schuecker et al. 2001). From the distribution of Fig.~2 one can compute the mean density of clusters as a function of distance, which indicates a very good completeness (i.e. constant density) out to at least $z=0.2$, possibly above. It is also evident how one can use the REFLEX survey to select sub-samples of clusters within well-defined ranges in $L_X$ (i.e. mass) and redshift. A number of follow-up studies of this kind are indeed ongoing (see B\\\"ohringer et al. 2001b for an overview), and more will certainly follow when the catalogue is released in early 2002. \\begin{figure} \\plotfiddle{guzzofig2.ps}{6.0cm}{0}{50}{50}{-160}{-180} \\caption{The distribution of X-ray luminosities (ROSAT band) as a function of redshift for all REFLEX clusters, with non-Abell clusters marked by filled circles. The two highest-redshift clusters are RXCJ1206.2-0848 at $z=0.441$ (whose RGB image is shown in Fig.~4) and RXCJ1347.4-1144, the most X-ray luminous cluster currently known, at $z=0.452$ (Schindler et al. 1995). } \\label{Lx_z} \\end{figure} The filled circles in Fig.~2 represent 142 clusters which do not appear in the visually-selected Abell/ACO catalogue (Abell 1958; Abell et al. 1989). It is interesting to see that these clusters are distributed basically at any redshift. This is a demonstration of how the ``richness'' criterion (the reason why most of these objects did not enter the Abell selection) is a poor indicator of the cluster mass, and how incomplete in mass a purely visually-selected cluster survey could be. We have also performed a direct X-ray analysis of the RASS data at all Abell-ACO cluster positions, measuring their X-ray flux. This gave the rather encouraging result that the REFLEX survey in fact detects {\\it all} Abell-ACO clusters within its flux and area boundaries\\footnote{Even more, the REFLEX survey misses only 1 of the so-called Supplementary Abell clusters i.e. the extension to the main ACO catalogue, where those objects that did not meet all original criteria while still looking as {\\it bonafide} clusters were listed.}. ", + "conclusions": "" + }, + "0112/astro-ph0112191_arXiv.txt": { + "abstract": "Despite many efforts to find a reasonable explanation, the origin of the \"knee\" in the cosmic ray spectrum at $E \\approx 10^{15.5} eV$ remains mysterious. In this letter we suggest that the \"knee\" may be due to a GZK-like effect of cosmic rays interacting with massive neutrinos in the galactic halo. Simple kinematics connects the location of the \"knee\" with the mass of the neutrinos, and, while the required interaction cross section is larger than that predicted by the Standard Model, it can be accommodated by a small neutrino magnetic dipole moment. The values for the neutrino parameters obtained from the analysis of existing experimental data are compatible with present laboratory bounds. ", + "introduction": " ", + "conclusions": "" + }, + "0112/astro-ph0112472_arXiv.txt": { + "abstract": "Radio variability on timescales from a few hours to several days in extragalactic flat-spectrum radio sources is generally classified as intra-day variability (IDV). The origin of this short term variability is still controversial and both extrinsic and intrinsic mechanisms must be considered and may both contribute to the observed variations. The measured linear and circular polarization of IDV sources constrains the low energy end of the electron population. Any population of cold electrons within sources at or above the equipartition temperature of $10^{11}$ K depolarizes the emission and can be ruled out. Intrinsic shock models are shown to either violate the large fraction of sources displaying IDV or they do not relax the light travel time argument for intrinsic variations. From structure function analysis, we further conclude that interstellar scintillation also leads to tight size estimates unless a very local cloud in the ISM is responsible for IDV. ", + "introduction": "Intraday Variability (IDV) of flat-spectrum compact Quasar cores and BL Lacs at cm-wavelength has been discovered in 1985 (Heeschen et al., 1987) and is a common phenomenon ($\\sim 25\\%$) (Quirrenbach et al., 1992) among these sources. On rare occasions in 0716+714 and 0954+658, correlations with optical variations have been found and the evidence has been reviewed several times (e.g. Wagner S.J., Witzel A., 1995; Krichbaum, this volume). Radio-optical correlations, if real, suggest either fast intrinsic variations or gravitational lensing. From the light travel time argument the apparent brightness temperatures $T_b$ are in the range of $10^{17}$--$ 10^{21}$K and far in excess of the intrinsic inverse Compton (IC) limit of $10^{12}$K. The observed superluminal motions of jet components in many of these sources imply $\\Gamma=5$--$10$ and allow for Doppler boosting factors ${\\cal{D}} =\\left(\\Gamma (1-\\beta \\cos\\theta)\\right)^{-1}$ and time shortening, which are insufficient for reducing the apparent $T_b$ down to the IC limit. The required Doppler factors $\\cal{D}$ are in the range of 60 up to 1000. Furthermore jets with a surface brightness at the IC limit are radiatively inefficient (Begelman, Rees \\& Sikora, 1994) and carry most of their energy as bulk motion. This is not supported by the observed power of radio lobes of these sources and raises the energy requirement for the central engine to an uncomfortable level. It is therefore argued (Readhead, 1994; Begelman, this volume) that incoherent synchrotron sources in jets should radiate at the equipartition temperature $T_E \\approx 10^{11}$\\,K. This enforces Doppler factors which are 2--3 times larger than for the IC limit. Furthermore synchrotron sources at the IC limit with bulk motions of $\\Gamma > 100$ are dominated by inverse Compton scattering of either AGN photons (Begelman, Rees \\& Sikora, 1994) or CMB photons at redshifts $z \\sim 1$ and not by the SSC process. The cooling is catastrophic, independent of the brightness temperature, and this explanation must therefore be discarded. Other more tempting suggestions are the propagation of relativistic thin shocks in the jets (Qian et al., 1991), so that the observed variability timescale is not a measure of the source size, and scintillation induced by the interstellar medium (e.g. Rickett 1990) of otherwise non-variable sources. We will explore both explanations in the following sections. ", + "conclusions": "Based on IDV observations of several Quasar and BL Lac radio cores with apparent brightness temperatures in the range of $10^{16}$--$10^{21}$\\,K we investigated the possibility of intrinsic variations due to the propagation of thin relativistic shocks. We find that no source with $T_b > 10^{16}$\\,K can be explained by that model without calling for Doppler factors larger than 20 and strong alignment between the jet and the line of sight. The required alignment firmly rules out this hypothesis as an explanation for all IDV sources. Based on refractive interstellar scintillation IDV can be explained by turbulence in the ISM. The scale size of the ionized gas responsible for IDV is about 100\\,pc in the case of 0917+624, and scintillation is quenched by the source size, which is larger than the refractive or Fresnel scale in the ISM. A fit to the structure functions at $\\lambda 20$cm and $\\lambda 6$cm indicates a steep power law ($\\beta \\approx 4.6$) for $\\Phi(q)$ fluctuations. This corresponds to an energy spectrum $\\sim q^{-2.6}$ that is much steeper than both a Kolmogorov spectrum and the $q^{-2}$ spectrum for compressible turbulence, but a bit shallower than $\\sim q^{-3}$ expected for two-dimensional turbulence. The length scale probed in the ISM by these measurements are between $5\\,10^8$\\,m and $2\\,10^9$\\,m at the peak of the structure functions. The slope of the turbulent spectrum is derived from structure functions at small timelags and the corresponding spectrum extends an order of magnitude down to smaller spatial scales. Compressible turbulence is not unexpected at these scales, if turbulence is driven by shocks from supernovae or by stellar winds. Scintillation in 0917+624 is quenched by the source size, which is one parameter of the theoretical fits to the SFs. The required sizes are $\\sim 40 \\mu$as and $\\sim 0.4$\\,mas at $\\lambda 6$cm and $\\lambda 20$cm respectively. At $\\lambda 6$cm, the flux of 0917+624 (redshift $z=1.446$) is 1.5Jy. Combined with the derived source size this implies $T_b = 10^{14}$K. Again Doppler factors of about 100 are needed to avoid the IC catastrophe or Doppler factors of 1000 to arrive at equipartition temperature. The sizes derived for 0917+624 from structure function models might be changed, if a degeneracy in the model parameters exists. The most plausible direction is an even closer ISM screen with a higher level of turbulence. In this case the angular Fresnel scale gets larger and larger source sizes are allowed." + }, + "0112/astro-ph0112466_arXiv.txt": { + "abstract": "In this {\\it Letter} we use a simple model to demonstrate the observational signatures of the magnetic connection between a black hole and a disk: (1) With the magnetic connection more energy is dissipated in and radiated away from regions close to the center of the disk; (2) The magnetic connection can produce a very steep emissivity compared to the standard accretion; (3) The observational spectral signature of the magnetic connection can be robust. These signatures may be identified with the observations of {\\it Chandra} and {\\it XMM-Newton}. In fact, the steep emissivity index for the Seyfert 1 galaxy MCG--6-30-15 inferred from the recent {\\it XMM-Newton} observation is very difficult to be explained with a standard accretion disk but can be easily explained with the magnetic connection between a black hole and a disk. ", + "introduction": "A magnetic field connecting a black hole to a disk has important effects on the balance and transfer of energy and angular momentum \\citep[and references therein]{bla99,bla00,li00a,li00b,li01,li00c}. When the black hole rotates faster than the disk, energy and angular momentum are extracted from the black hole and transferred to the disk. Thus, the rotational energy of the black hole provides an energy source for the radiation of the disk in addition to disk accretion. The energy deposited in the disk by the black hole can be dissipated by the internal viscosity of the disk and radiated away to infinity \\citep{li00a,li00b}. A disk powered by a black hole through magnetic connection shows interesting features different from that of a standard accretion disk \\citep{li00b,li01}. The energy radiated by the disk comes from regions closer to the center of the disk, the radiation flux decreases more rapidly with radius, which approaches $r^{-3.5}$ at large radii (compared to $r^{-3}$ for a standard accretion disk). This implies that the radiation spectrum of a disk with magnetic connection observed by a distant observer will also be different from that of a standard accretion disk. Interestingly, the most recent {\\it XMM-Newton} observation of the nearby bright Seyfert 1 galaxy MCG--6-30-15 reveals an extremely broad and red-shifted Fe K$\\alpha$ line indicating its origin from the very most central regions of the accretion disk \\citep{wil01}. To explain the observed spectrum a very steep emissivity profile with index $\\alpha = 4.3-5.0$ is required [a steep emissivity for the same galaxy and Mrk-766 is also reported by \\citet{bra01}]. Such a steep emissivity is very difficult to be explained within the framework of a standard accretion disk, but the magnetic connection between a black hole and a disk may provide a natural explanation \\citep{wil01}. Though there is yet another possible explanation in term of the magnetic connection between the inner boundary of the disk and the plunging material inside the inner boundary \\citep{kro99,gam99,ago00}, the issue remains controversial \\citep{pac00} and the most recent numerical simulations show that the stress produced by such a magnetic connection sensitively depends on the thickness of the disk: the stress in the plunging region is significantly reduced as the thickness of the disk decreases \\citep{haw01b,haw01c,arm01,haw01a}\\footnote{In fact, all analytic and numerical works on disk dynamo show that the strength of the magnetic field lying in the disk is always limited by the gas pressure in the disk \\citep{gal79,bal98,mil00}. Since in the plunging region the gas pressure is very low due to the supersonic flow we expect that the stress produced by the magnetic field in the plunging region cannot be important.}. In this {\\it Letter} we use a simple model to demonstrate the observational signatures of the magnetic connection between a black hole and a disk. We assume that an axisymmetric magnetic field connects a black hole to a non-accretion disk from the inner boundary of the disk to a circle with radius $r_b$. As in the case of a standard accretion disk, the inner boundary of the disk is assumed to be at the marginally stable orbit with radius $r_{ms}$. We will calculate the radiation flux, the emissivity index, and the radiation spectrum observed by a distant observer for various distribution of the magnetic field between $r_{ms}$ and $r_b$ on the disk. We will compare the results with that of a standard accretion disk and look for the robustness of the observational signatures of the magnetic connection. ", + "conclusions": "Ever since \\citet{pen69} proposed the first Gedankenexperiment for extracting energy from a rotating black hole many people have looked for more practical ways that may work in astronomy. Among many alternatives the Blandford-Znajek mechanism \\citep{bla77,mac82,phi83} has been thought to be promising for powering extragalactic jets. The magnetic connection between a rapidly rotating black hole and a disk is a variant of the Blandford-Znajek mechanism and is more efficient in extracting energy from the black hole \\citep{li00a}. The energy extracted from the black hole is deposited into the disk, dissipated by the internal viscosity of the disk, and subsequently radiated away to infinity \\citep{li00b}. Such a disk can radiate without accretion, the power of the disk comes from the rotational energy of the black hole. The magnetic connection not only increases the radiation efficiency of the disk (i.e., increases the ratio of radiated energy to accreted mass) but also produces observational signatures. With the magnetic connection, more energy is dissipated in and radiated away from regions closer to the center of the disk, which in turn produces a steep emissivity profile. The spectral signature produced by the magnetic connection is also significantly different from that produced by the standard accretion. The recent {\\it XMM-Newton} observations of soft X-ray emission lines from two Narrow Line Seyfert 1 galaxies MCG --6-30-15 and Mrk 766 shows an extreme steep emissivity profile, which has been suggested to indicate that most of the line emission originates from the inner part of a relativistic accretion disk \\citep{bra01,wil01}. From Fig.~\\ref{fig1} and Fig.~\\ref{fig2} we see that these features are just predicted by the magnetic connection between a rapidly rotating black hole and a disk. Another kind of observation which may be relevant to the magnetic connection between a black hole and a disk is the kilohertz quasi-periodic oscillations (kHz QPOs) in X-ray binaries, which has been suggested to originate from the inner edge of a relativistic accretion disk \\citep[and references therein]{van00}. The magnetic connection between a Kerr black hole and a disk provides an interesting model for kHz QPOs. In the case of strong coupling, \\cite{bla99} has proposed to associate QPOs with the radial oscillation produced by the magnetic connection. There is yet another possibility: QPOs may be produced by a non-axisymmetric magnetic field connecting a black hole to a disk \\citep{li00b,li01}. Suppose a bunch of magnetic field lines connect a black hole to a disk and the feet of the magnetic field lines are concentrated in a small region in the disk, then a hot spot will be produced on the disk surface if the black hole rotates faster than the disk. Since the disk is perfectly conducting, the magnetic field is frozen in the disk so the hot spot will co-rotate with the disk and show a periodic oscillation. While the model presented in this {\\it Letter} is so simple that it cannot be directly applied to comparison with observations, it demonstrates some interesting observational signatures of the magnetic connection between a black hole and a disk. With {\\it Chandra} and {\\it XMM-Newton} telescopes it becomes possible to probe the inner region of a disk around a black hole, so the observational signatures of the magnetic connection may be identified." + }, + "0112/astro-ph0112185_arXiv.txt": { + "abstract": "We describe a practical method for constructing axisymmetric two-integral galaxy models (with distribution functions of the form $f(E,L_z)$, in which $E$ is the orbital energy, and $L_z$ is the vertical component of the angular momentum), based on Schwarzschild's orbit superposition method. Other $f(E,L_z)$-methods are mostly based on solving the Jeans equations or on finding the distribution function directly from the density, which often places restrictions on the shape of the galaxy. Here, no assumptions are made and any axisymmetric density distribution is possible. The observables are calculated (semi-)analytically, so that our method is faster than most previous, fully numerical implementations. Various aspects are tested extensively, the results of which apply directly to three-integral Schwarzschild methods. We show that a given distribution function can be reproduced with high accuracy and investigate the behaviour of the parameter that is used to measure the goodness-of-fit. Furthermore, we show that the method correctly identifies the range of cusp clopes for which axisymmetric two-integral models with a central black hole do not exist. ", + "introduction": "\\label{sec:intro} The fundamental quantity in galaxy dynamics is the distribution function (DF), which specifies the distribution of the stars in a galaxy over position and velocity. By Jeans' theorem, the DF is a function of the isolating integrals that are conserved by the corresponding potential (e.g., Lynden-Bell 1962; Binney 1982). Axisymmetric galaxies, which we will study here, conserve at least the two classical integrals of motion, the energy $E$ and the vertical component of the angular momentum $L_z$. The DF can generally not be measured directly, but since observationally accessible quantities, such as the projected density and the line-of-sight velocity, are simple moments of the DF, photometric and kinematic observations can provide information on its properties. In some cases, the part of a two-integral distribution function that is even in the velocities, $f(E,L^2_z)$, can be obtained analytically by using integral transforms to solve the relation between the DF and the intrinsic density $\\rho(R,z)$, where $(R,\\phi,z)$ are the usual cylindrical coordinates. The odd part can be found similarly from $\\rho \\langle v_\\phi \\rangle$, where $\\langle v_\\phi \\rangle$ is the mean streaming velocity. These Laplace (Lynden-Bell 1962; Lake 1981), Stieltjes (Hunter 1975b) and Laplace-Mellin (Dejonghe 1986) transforms have the drawback that numerical implementation is difficult and that they require $\\rho$ and $v_\\phi$ to be explicit functions of the potential $V$ and the cylindrical radius $R$. A more general formalism (Hunter \\& Qian 1993; Qian et al.\\ 1995, hereafter Q95), the HQ-method, uses contour integration instead of integral transforms. This means it is simpler to implement, does not explicitly require $\\rho(R,V)$ and $v_\\phi(R,V)$, but a suitable contour for the integration has to be chosen, which may not always be at hand. Because of the drawbacks of these analytical formalisms, numerical methods that can be applied to arbitrary potential-density pairs are more attractive. Various methods have been developed, many of which circumvent knowledge of the DF by solving the Jeans equations directly (van der Marel et al.\\ 1994; Magorrian 1995), while others assume that the DF can be represented by a superposition of basis functions (Dehnen \\& Gerhard 1994; Kuijken 1995). Accurate estimates of the mass-to-light ratio $M/L$ can be obtained in both ways (van der Marel 1991; Shaw et al.\\ 1993; van den Bosch 1996). The predicted central black hole masses that are obtained with two-integral models (Magorrian et al.\\ 1998) seem to overestimate the true values, but are still very useful to narrow down the parameter range for more sophisticated modeling (van der Marel et al.\\ 1998, hereafter vdM98; Gebhardt et al.\\ 2000; Bower et al.\\ 2001). Numerical orbit integrations and observations of, e.g., the anisotropy of the stellar dispersions in the solar neighbourhood show that, in addition to $E$ and $L_z$, a third integral is conserved for most orbits. In separable potentials, this third integral is exact and has a closed form (e.g., Kuzmin 1956; de Zeeuw 1985), allowing a direct calculation of the DF (e.g., Dejonghe \\& de Zeeuw 1988). The most general family of these potentials corresponds to flattened mass models with constant-density cores (de Zeeuw, Franx \\& Peletier 1986), which provide a poor description of the inner regions of most galaxies, since these contain a central density cusp (Lauer et al.\\ 1995). We conclude that many of the existing two-integral, as well as the analytical three-integral methods, are applicable to a limited range of galaxy models. A flexible method for calculating numerical galaxy models was designed by Schwarzschild (1979, 1982), who represented the DF numerically by the occupation numbers in a superposition of building blocks. There are no restrictions on the density or the potential, and no a priori assumptions have to be made about the shape or the degree of anisotropy of the galaxy. Schwarzschild's original implementation was aimed at reproducing a given triaxial density distribution and was subsequently applied to a large variety of galaxy models, from spherically and axially symmetric (Richstone \\& Tremaine 1984; Levison \\& Richstone 1985) to triaxial density distributions (e.g., Vietri 1986; Statler 1987; Schwarzschild 1993, Merritt \\& Fridman 1996; Siopis \\& Kandrup 2000). More general versions of the method can also reproduce kinematical observations of spherically and axially symmetric galaxies that obey up to three integrals of motion (Zhao 1996; Rix et al.\\ 1997, hereafter R97; Cretton et al.\\ 1999, hereafter C99; Cretton et al.\\ 2000; H\\\"afner et al.\\ 2000).\\looseness=-2 In most implementations of Schwarzschild's method, orbits are used as individual building blocks. In non-separable potentials the third integral of motion is not known explicitly, so that the orbital properties can only be obtained by solving the equations of motion numerically. This orbit integration has to be carried out until the relevant quantities have averaged out and do not change upon a new time-step (Pfenniger 1984; Copin et al.\\ 2000). Especially for orbits that are near-stochastic, or that have nearly commensurate fundamental frequencies, this condition is difficult to achieve and the orbit integration can be very time-consuming. Furthermore, an orbit fills a region in space that can have sharp edges, depending on the combination of integrals of motion that it obeys. A superposition of a finite (and relatively small) number of orbits can therefore show artefacts that are caused by these edges. One way to solve this is to `blur' orbits in phase-space randomly, another is to use building blocks that are smoother (e.g., Zhao 1996; Merritt \\& Fridman 1996; R97; C99). Of particular interest are the so-called isotropic and two-integral components (ICs and TICs). Isotropic components, which are completely specified by their energy $E$, fill the volume inside the equipotential at $E$, while TICs, which are fixed once $E$ and $L_z$ are chosen, completely fill the corresponding zero-velocity curve (ZVC). The advantages of these building blocks is apparent: both the equipotential and the ZVC are smooth surfaces, and ICs and TICs can be considered as weighted combinations of all orbits with the same energy $E$ (or $E$ and $L_z$), including the irregular orbits. These building blocks can be used in Schwarz\\-schild's method in two ways. They can be added to every energy or every combination $(E,L_z)$ of a (three-integral) orbit library, which will `automatically' take care of the stochastic orbits (Zhao 1996; R97; C99). Alternatively, a fully isotropic or two-integral model can be constructed by using only ICs or TICs as building blocks. C99 suggested that this might be a practical way of constructing such models, and derived some analytic properties, but did not pursue these models further. We do so in this paper and concentrate on the two-integral case, since the properties of the isotropic components follow from those of the TICs by taking $L_z=0$. In \\S\\ref{sec:method}, we collect the properties of the TICS, and summarize how we use them in Schwarzschild's method. The numerical aspects are discussed in detail in \\S\\ref{sec:implementation}. We show in \\S\\ref{sec:examples} that our method can reproduce a model with a known analytical distribution function with high accuracy. We also study a set of mass models with a central black hole, and show that our method is able to detect when a self-consistent solution does not exist. We summarize our conclusions in \\S\\ref{sec:summary}. Applications of our software to model the observed two-dimensional surface brightness and kinematics of nearby elliptical galaxies observed with {\\tt SAURON} and {\\tt STIS} are described in two follow-up papers (Verolme et al.\\ 2002, in preparation; McDermid et al.\\ 2002, in preparation). ", + "conclusions": "\\label{sec:summary} We have presented an alternative to the existing analytical and numerical methods for calculating two-integral distribution functions. It is based upon Schwarzschild's orbit superposition method and able to deal with arbitrary density profiles and galaxy potentials. Instead of orbital building blocks, we use the so-called two-integral components, which are smoother and implicitly include stochastic orbits. Due to their delta-function behaviour, the observables can be calculated (semi-)analytically, which speeds up the calculations considerably. We checked that the method is able to reproduce a known combination of potential, density and distribution function and tested the regularisation algorithm in the process. This test shows that fitting the meridional plane masses alone is not enough to constrain the DF, while including a (modest) amount of regularisation is enough to smoothen the DF toward the theoretical curve. We have also tested the $\\chi^2$ (\\ref{eq:chisq}), which measures the quality-of-fit of the resulting model, against analytical investigations. This demonstrates that a high value of $\\chi^2$ can be taken as a sure indication that the model does not exist. These results show that our method is flexible and reproduces analytical results very closely. Additionally, many of the characteristics of the implementation that we used, such as the regularisation method and the $\\chi^2$-parameter, are used in more general (three-integral) models. Regularisation is therefore also a very useful and necessary tool to smoothen the distribution function that is found by these methods, and the $\\chi^2$-parameter is a useful diagnostic to test the existence of the resulting models. Applications in which we calculate fully three-integral models (including TICs) of observed kinematics will follow in a subsequent paper. \\medskip \\noindent{\\bf Acknowledgements}\\\\ \\noindent It is a pleasure to thank Glenn van de Ven, Michele Cappellari and Roeland van der Marel for useful discussions and a critical reading of the manuscript. Furthermore, we are very thankful to Roeland van der Marel for kindly providing his three-integral Schwarzschild software." + }, + "0112/astro-ph0112520_arXiv.txt": { + "abstract": "We present the results of UV imaging polarimetry of the Seyfert 2 galaxy Mrk 477 taken by the Faint Object Camera onboard the Hubble Space Telescope (HST). From a previous HST UV image ($\\lambda \\sim 2180$\\AA), Mrk 477 has been known to have a pointlike bright UV hotspot in the central region, peculiar among nearby Seyfert 2 galaxies. There are also claims of UV/optical variability, unusual for a Seyfert 2 galaxy. Our data show that there is an off-nuclear scattering region $\\sim 0.''6$ ($\\sim 500$ pc) NE from the hotspot. The data, after the subtraction of the instrumental effect due to this bright hotspot region, might indicate that the scattered light is also detected in the central $0''.2$ radius region and is extended to a very wide angle. The hotspot location is consistent with the symmetry center of the PA pattern, which represents the location of the hidden nucleus, but our data do not provide a strong upper limit to the distance between the symmetry center and the hotspot. We have obtained high spatial resolution color map of the continuum which shows that the nuclear spiral arm of $0.''4$ scale ($\\sim$ 300pc) is significantly bluer than the off-nuclear mirror and the hotspot region. The nature of the hotspot is briefly discussed. ", + "introduction": "Mrk 477 is classified as a Seyfert 2 galaxy since its permitted emission lines are narrow in the optical spectrum. However, it has some unusual properties among this class of objects. Seyfert 2 galaxies generally do not show variability in the UV/optical, but for this object, some variations have been reported. De Robertis (1987) first reported that the Fe lines ([Fe VII] $\\lambda$6087 and [Fe X] $\\lambda$6375) are variable, and the optical continuum seems to have increased by a factor of 2 between the 1980 and 1985 observations. Although Veilleux (1988) pointed out that the reported variation in the FeX line is due to the afterglow in the detector and the FeVII line variation is questionable due to the presence of a sky line, Kinney et al. (1991) also reported that the UV/optical continuum increased by a factor of 2 over a $5-6$ year period, based on the mismatch of the flux at $\\sim 3000$\\AA\\ from the IUE data taken in 1984 (partly in 1983) and the optical spectra taken in 1989 (confirmed in 1990). Kinney et al. also suggested that its nuclear ionizing source does not need to be blocked from direct view based on a photon budget argument. This is strange in the sense that, if we are seeing the nucleus directly, we should see broad permitted lines also directly, since the nuclear continuum source is generally thought to be more compact than the broad-line region. But in the Mrk 477 spectrum, we do not see the broad lines in total flux (see discussion of upper limit in \\S\\ref{sec-disc}). This might lead us to consider an actual lack of a broad-line region in this object. However, optical spectropolarimetry has shown the presence of broad lines in the polarized flux spectrum (Tran, Miller, \\& Kay 1992; Tran 1995), suggesting that the broad-line region does exist but is hidden from direct view and seen only through scattered light. This is in accordance with the general idea for Seyfert 2 galaxies that the continuum source and broad-line region seen in Seyfert 1 galaxies exist also in Seyfert 2 galaxies but are hidden from direct view (Antonucci \\& Miller 1985). This idea is consistent with the fact that nearby Seyfert 2 galaxies generally do not show an unresolved bright nuclear source in their HST UV/optical image (Nelson et al. 1996; Malkan, Gorjian, \\& Tam 1998). However, the HST UV image (at $\\sim 2180$\\AA; Heckman et al. 1997) has shown that Mrk477 has a fairly bright UV pointlike source in its central region, which is again peculiar. On the other hand, the HST/GHRS UV spectroscopy has shown that there is strong starburst activity in a rather compact nuclear region (a few hundred pc scale; Heckman et al. 1997). In order to reveal the detailed nuclear structure of this peculiar Seyfert 2 galaxy, we have conducted HST imaging polarimetry of Mrk 477. Our data have spatially resolved the nuclear polarization structure. We describe our observations in \\S\\ref{sec-obs} and the results in \\S\\ref{sec-res}. The implication of the results are discussed in \\S\\ref{sec-disc} and our conclusions are summarized in \\S\\ref{sec-conc}. We adopt $H_0 = 65$ km sec$^{-1}$ Mpc$^{-1}$ throughout this paper. Mrk 477 is at $z=0.038$, so the distance is $\\sim$ 180 Mpc and one arcsec corresponds to $\\sim$ 800 pc. The Galactic reddening for this object is low, $E(B-V) = 0.011$ (NED; Schlegel et al. 1998). For the Galactic reddening correction, we adopt the reddening curve of Cardelli et al. (1989) with $A_V/E(B-V) = 3.1$. ", + "conclusions": "\\label{sec-conc} We have presented HST UV imaging polarimetry data of the Seyfert 2 galaxy Mrk 477. For this galaxy, there are claims of variability in the UV/optical, unusual for a Seyfert 2 galaxy. It has a UV bright pointlike hotspot in the central region, which is also peculiar among nearby Seyfert 2 galaxies. Our data identify an off-nuclear scattering region $\\sim 0.''6$ ($\\sim 500$ pc) NE from the hotspot. The data, after the subtraction of the instrumental effect from the bright hotspot region, might indicate that the scattered light is also detected in the nuclear vicinity ($\\sim 0.''2$ radius) and is extended widely with a full opening angle of $\\sim 180$\\degr\\ around the hotspot region. This could lead us to consider the possibility that our line of sight is grazing the matter obscuring the nucleus, which might be the cause of some of the peculiar properties of this galaxy. However, the uncertainty from the subtraction process is large and we need more evidence to support this claim. The hotspot location is consistent with the symmetry center of the PA pattern, which represents the location of the hidden nucleus, but our data do not provide a strong upper limit on the distance between the hotspot and the symmetry center. The hotspot radiation seems to be slightly polarized, but it does not appear to be dominated by scattered light. Since we essentially do not see direct broad lines and the UV spectral shape of the hotspot is not too red, we do not expect the hotspot to be a heavily absorbed nucleus which is usually thought to be inside the broad-line region. There would be another continuum source outside the broad-line region, or the nuclear continuum source has to be larger than the broad-line region." + }, + "0112/hep-ph0112153_arXiv.txt": { + "abstract": "Strange quark matter in a color-flavor locked state is significantly more bound than ``ordinary'' strange quark matter. This increases the likelihood of strangelet metastability or even absolute stability. Properties of color-flavor locked strangelets are discussed and compared to ordinary strangelets. Apart from differences in binding energy, the main difference is related to the charge. A statistical sample of strangelets may allow experimental distinction of the two. Preliminary estimates indicate that the flux of strangelets in galactic cosmic rays could be sufficient to allow for strangelet discovery and study in the upcoming Alpha Magnetic Spectrometer AMS-02 cosmic ray experiment on the International Space Station. ", + "introduction": "It has long been known that phenomenological strong interaction models, most notably the MIT bag model, allow absolute stability (energy per baryon below 930~MeV) of three flavor quark matter for certain ranges of parameters, and metastability for a wider parameter span. While bag model calculations are clearly only a crude approximation to full (but untractable) QCD, the confirmation of strange matter (meta)stability would have important consequences. For instance strange matter stability would imply that ``neutron stars'' are actually quark stars (strange stars), and metastability could also significantly change the physics of neutron star interiors. Even if quark matter is (meta)stable in bulk, finite size effects (surface and curvature energies) increase the energy per baryon for small lumps of three flavor quark matter, called strangelets. This makes it less likely to form such objects in heavy-ion collisions, and such formation is also hindered by the high entropy/temperature environment, which destabilizes strangelets further (``making ice-cubes in a furnace''). Nevertheless, tiny amounts of strangelets might be formed in colliders, and the detection of one would be the ultimate smoking gun for the quark-gluon plasma \\cite{madsen99}. The probability of strangelet (meta)stability is increased with the recent demonstration that quark matter at high density may be in a so-called color-flavor locked phase where quarks with different color and flavor quantum numbers form Cooper pairs with pairing energy $\\Delta$ perhaps as high as 100~MeV \\cite{alford01b}. Such a state is significantly more bound than ordinary quark matter, and this increases the likelihood that quark matter composed of up, down, and strange quarks may be metastable or even absolutely stable. In other words color-flavor locked quark matter rather than nuclear matter could be the ground state of hadronic matter. In the following I shall briefly summarize the physics of ``ordinary'' strangelets, followed by a summary of recent work on color-flavor locked strangelets. Finally, I discuss the potential of experimental discovery of strangelets with the AMS-02 detector on the International Space Station, which may even (by studying the mass-charge relation) allow a distinction between ordinary and color-flavor locked strangelets. \\begin{figure}[h!] \\begin{center} \\includegraphics[height=10cm]{fig1.ps} \\caption{Energy per baryon in MeV is plotted as a function of baryon number for ``ordinary'' strangelets with s-quark mass of 50, 150, and 300~MeV (bottom to top). The bag constant is $B=(145~{\\rm MeV})^4$. Mode filling within the MIT bag model corresponds to the solid curves, whereas dashed curves were calculated from the smoothed density of states within the multiple reflection expansion. It is clearly seen how $E/A$ grows from a bulk value (c.f.~Fig.~2) at high $A$ to a value significantly increased by finite size effects at low $A$. It is also seen how the multiple reflection expansion including volume, surface, and curvature terms (dashed curves) gives an excellent reproduction of the overall behavior of $E/A$. The upper set of curves (for $m_s=300$~MeV) essentially corresponds to two flavor quark matter, since only up- and down-quarks are energetically favorable for such a high strange quark mass. The bag constant chosen is close to the lower bound permitted if nuclei should remain stable against direct decay into two flavor quark matter (decay into strangelets is forbidden since it requires a high order weak interaction to create strange quarks). } \\label{fig:ordinary} \\end{center} \\end{figure} ", + "conclusions": "" + }, + "0112/astro-ph0112070_arXiv.txt": { + "abstract": "{ A homogeneous set of $UBV$ photometry (354 data points obtained between 1983 and 1998) for the Be/X-ray binary \\object{A0535+26}~= \\object{V725~Tau} is analysed, aiming to look for possible periodic component(s). After subtraction of the long-term variation it was found that only a $\\approx 103^d$ periodic component remains in the power spectra in both the $V$ and $B$ colour bands. The probability of chance occurrence of such a peak is less than 0.1\\%. There are no signs of optical variability at the X-ray period ($\\approx 111^d$). We discuss possible reasons for a 103-day modulation and suggest that it corresponds to a beat frequency of the orbital period of the neutron star and the precession period ($\\approx 1400^d$) either of an accretion disc around the neutron star or a warped decretion disc around the Be star. ", + "introduction": "The X-ray binary \\object{A0535+26} has been an object of interest to both observers and theorists for a quarter of a century, from the moment of its first documented X-ray outburst in 1975 (Rosenberg et al.~\\cite{Ros75}). As soon as an optical counterpart was identified almost simultaneously by several authors (Liller~\\cite{Lil75}; Murdin~\\cite{Mur75}) with the 9th magnitude Be star \\object{HDE~245770}, later named \\object{V725~Tau}, there were several attempts to find the orbital period of the system, using both X-ray and optical (spectral and photometric) data. Priedhorsky \\& Terrell (\\cite{PT83}) have shown that all but the most powerful X-ray outbursts (1975 April and 1980 October) take place with a period $\\approx$111 days. Now the general opinion is that the orbital period is 110--111$^d$. The most reliable determination of the X-ray period and initial epoch is currently that of Motch et al. (\\cite{Mo91}): $\\mbox{P}_\\mathrm{XR} = 111\\fd 38\\pm 0\\fd 11$ and $T_0 = \\mbox{JD} 2446734.3\\pm 2\\fd 6$. This and other periodicities were claimed to have been found in optical (photometric and spectral) data. For instance, Hutchings et al. (\\cite{Hu78}) found that the radial velocities of absorption lines of \\object{HDE~245770} are variable, and suggested several probable periods: 28.6, 48 and 94 days. Some years later and based on a very similar dataset, Hutchings (\\cite{Hu84}) found a 112 day period. Guarnieri et al. (\\cite{Gu82}) noted at first possible 32, 63 or 77 day periods in their photometric data, but later (Guarnieri et al. \\cite{Gu85}) found modulation in the $V$ band photometry with the proposed orbital period of 110 days. Gnedin et al. (\\cite{Gn88}) made a Fourier analysis of $V$ band photometry obtained during 1981--1985 and marked out periods of 1100, 103 and 28 days, but found no periodicity corresponding to the X-ray period. Besides the determinations noted above, there are several more publications on this subject, listed by Giovannelli and Graziati (\\cite{GG92}). In a recent paper by Hao et al. (\\cite{Hao96}) all the published data sets from 1981 to 1993 have been analysed together. Their analysis have shown only two significant periods -- 830 and 507 days. However their light curve modeling does not agree with the observational data and their prognosis of the photometric behavior of \\object{HDE~245570} for 1994--1997 was not confirmed by later observations (see Fig.~1 in Hao et al. \\cite{Hao96}, Fig.~1 in Lyuty \\& Zaitseva \\cite{LZ2000} and Fig.~\\ref{fig:2} in this work). It should be noted also that in Hao et al.~(\\cite{Hao96}) no correction has been made for different zero-points and colour systems of the data sampled by different groups. As soon as our data set had grown substantially and surpassed that used in Gnedin et al. (\\cite{Gn88}), as well as other published data sets, we attempted to search for periodic components in the light curves of \\object{A0535+26}/\\object{V725~Tau}. In Sect.~2 we describe our techniques of data analysis and obtain a value of the optical period, distinctly different from the orbital one, while in Sect.~3 we discuss possible reasons of the light modulation and suggest that it is caused by interplay of the orbital and precession motions. In Sect.~4 we summarize the results. ", + "conclusions": "An analysis of the uniform photometric data set obtained in the period 1983--1998 has allowed a confident separation of the periodic constituent in the light curve of the high-mass X-ray binary, A0535+26/V725~Tau. The parameters of this periodic component and its link with the phase of activity of the optical component allow us to suggest precession of an accretion disc around the neutron star or a warped equatorial disc around a Be star as the most likely mechanisms. At this point, both models seem to be viable; meanwhile, the analysis of already existing spectral data could be helpful, in the sense that if the warped disc model does reflect reality, one might expect to see $V/R$ and $EW$ variations corresponding to the precessional motion. Within both models, we do not expect that substantial X-ray outbursts to occur during ascending parts of the large-scale optical light curve. Moreover, X-ray outbursts tend to occur at specific phases of the $102\\fd 8$ optical light curve. Taken together, both effects can explain the ``missing outburst'' phenomenon. Our results lead us to suppose that in other similar systems we can hope to distinguish the minor photometric variations close to, but not necessarily coincident with, the binary rotation period. In future papers of this series, we plan to apply the techniques described here to the analysis of the light curves of similar systems, such as \\object{X~Per} and \\object{V635~Cas}." + }, + "0112/astro-ph0112300_arXiv.txt": { + "abstract": "SETI@home observes a 2.5 MHz bandwidth centered on 1420 MHz near the 21-cm line using a short line feed at Arecibo which provides a 6$^\\prime$ beam. This feed sits on Carriage House 1. During normal astronomical observations with the new Gregorian dome the feed scans across the sky at twice the sidereal rate. We are using the SETI@home receiver to obtain about $4.4\\times 10^{6}$ H{\\sc i} spectra per year with integration time of 5 seconds per spectrum. We have accumulated 2.6 years of data covering most of the sky observable from Arecibo. This survey has much better angular resolution than previous single dish surveys and better sensitivity than existing or planned interferometric surveys. ", + "introduction": " ", + "conclusions": "" + }, + "0112/astro-ph0112136_arXiv.txt": { + "abstract": "We compute spectra from accretion disks around rapidly rotating neutron stars. The full effect of general relativity is considered for the structure calculation of the stars. We take into account the Doppler shift, gravitational redshift and light-bending effects in order to compute the observed spectra. To facilitate direct comparison with observations, a simple empirical function is presented which describes the numerically computed spectra well. This function can in principle be used to distinguish between the Newtonian spectra and the relativistic spectra. We also discuss the possibility of constraining neutron star's equation of state using our spectral models. ", + "introduction": "A large number of low mass X-ray binaries (LMXB) are believed to harbor neutron stars, rotating rapidly due to accretion-induced angular momentum transfer. These systems show many complex spectral and temporal behaviors. One of the main purpose for studying such behaviors is to understand the properties of very high density $(\\sim 10^{15}$ g cm$^{-3})$ matter at the neutron star core. Such high densities can not be created in the laboratory and only the study of these sources can give a possible answer to this fundamental question of physics. Here we calculate the equation of state (EOS) dependent model spectra of the accretion disks around rapidly rotating neutron stars. These models, when fitted to the observed spectra, can in principle help to constrain EOS models and hence to understand the properties of core-matter of neutron stars. ", + "conclusions": "" + }, + "0112/astro-ph0112171_arXiv.txt": { + "abstract": "In this talk I will present a model for primordial galaxy formation. In particular, I will review the feedback effects that regulate the process: (i) radiative (i.e. ionizing and H$_2$--photodissociating photons) and (ii) stellar (i.e. SN explosions) feedback produced by massive stars. I will show the results of a model for galaxy formation and IGM reionization, which includes a self-consistent treatment of the above feedback effects. Finally, I will describe a Monte Carlo method for the radiative transfer of ionizing photons through the IGM and discuss its application to the IGM reionization problem. ", + "introduction": "The application of the Gunn-Peterson test to QSOs absorption spectra suggests that the Intergalactic Medium (IGM) is completely reionized by $z \\sim 6$. Several authors claim that the known population of quasars and galaxies provides $\\sim 10$ times fewer ionizing photons than are necessary to keep the observed IGM ionization level; thus, additional sources of ionizing photons are required at high redshift, the most promising being early galaxies and quasars. Recent observational evidences suggest the existence of an early population of pregalactic stellar objects which could have contributed to the reionization and metal enrichment of the IGM. In the following section, I will review the processes that regulate the formation and the evolution of such galaxies. ", + "conclusions": "To properly describe the IGM reionization process, the following ingredients are needed: i) a reliable model for structure formation, including a self-consistent treatment of feedback effects and the best possible balance between large volume and high resolution of the simulation; ii) an accurate treatment of the radiative transfer of ionizing photons; iii) a better constraint on the parameters involved in the calculation, such as the source emission properties and the escape fraction." + }, + "0112/astro-ph0112492_arXiv.txt": { + "abstract": "We present moderate (1$''$) and high resolution (0.2$''$) observations of the CO (2--1) emission at 43 GHz, and radio continuum emission at 1.47 GHz, from the $z = 4.7$ QSO BRI 1202--0725 and the $z = 4.4$ QSO BRI~1335--0417 using the Very Large Array. The moderate resolution observations show that in both cases the CO emission is spatially resolved into two components separated by 1$''$ for 1335--0417 and 4$''$ for 1202--0725. The high resolution observations show that each component has sub-structure on scales $\\sim 0.2''$ to 0.5$''$, with intrinsic brightness temperatures $\\ge 20$ K. The CO ladder from (2-1) up to (7-6) suggests a high kinetic temperature for the gas ($\\rm T_{kin} \\simeq 70$ K), and a high column density (N(H$_2$) $\\simeq 10^{24}$ cm$^{-2}$). In both sources the continuum-to-line ratio: ${{L_{\\rm FIR}}\\over{L'_{\\rm CO(1-0)}}} \\simeq 335$. All these characteristics (brightness temperature, excitation temperature, column density, and continuum-to-line ratio) are comparable to conditions found in low redshift, ultra-luminous nuclear starburst galaxies. We find that the CO emitting regions in 1202--0725 and 1335--0417 must be close to face-on in order to avoid having the gas mass exceed the gravitational mass, implying perhaps unreasonably large rotational velocities. While this problem is mitigated by lowering the CO luminosity-to-H$_2$ mass conversion factor (X), the required X values become comparable to, or lower than, the minimum values dictated by optically thin CO emission. We considered the possibility of magnification by gravitational lensing in order to reduce the molecular gas masses. ", + "introduction": "By selecting for very red, point-like optical sources, McMahon (1991) has identified a large sample of $z \\ge 4$ QSOs from the Automatic Plate Measuring survey (Irwin, McMahon, \\&\\ Hazard 1991), including the two sources BRI 1202--0725 at $z = 4.7$ and BRI 1335--0415 at $z = 4.4$. Optical spectra of both of these sources show strong associated Ly$\\alpha$ absorption (Storrie-Lombardi et al. 1996). \\subsection{1202--0725} Omont et al. (1996) and Ohta et al. (1996) detect high order CO line emission from 1202--0725, as well as thermal continuum emission from warm dust at 1.35mm. The continuum and line emission consist of a double source with an angular separation of 4$''$. The total flux density of the source at 1.2mm is $12 \\pm 3$ mJy, implying a far IR luminosity of $L_{\\rm FIR} = 4.2 \\times 10^{13}$ L$_\\odot$, assuming a spectral energy distribution (SED) typical for an ultra-luminous infrared galaxy (ULIRG; $L_{\\rm FIR} \\sim 10^{12}$), where $L_{\\rm FIR}$ is defined as in Helou et al. (1988) and Condon (1992), ie. the integrated luminosity between rest frame wavelengths of 42 to 122 $\\mu$m. The southern source in 1202--0725 comprises about 65$\\%$ of the total. Observations of this source with ISO at mid- to far-IR wavelengths constrain the dust temperature to be $\\simeq 70$ K for a dust emissivity index $\\beta = 1.5$ (Leech, Metcalfe, \\& Altieri 2001). The CO observations of Omont et al. (1996a) and Guelin et al. (2001) indicate a difference in the CO line widths for the northern and southern components. Gaussian fitting to the CO(5-4) line emission profiles by Omont et al. (1996) results in $z = 4.6947$, line Full Width at Half Maximum (FWHM) = 190 km s$^{-1}$, and an integrated line flux density = 1.1$\\pm$0.2 Jy km s$^{-1}$ for the southern source, and $z = 4.6916$, FWHM = 350 km s$^{-1}$, and an integrated line flux density = $1.3 \\pm 0.3$ Jy km s$^{-1}$ for the northern source. This difference is important, since it argues against gravitational lensing as the origin for the observed double source and large apparent luminosity. However, Guilloteau (2001) has recently called the difference in CO profiles into question due to the possible contamination of the line emission from the northern source by continuum emission at 3mm. The data presented herein support the original conclusions of Omont et al. (1996a) and Guelin et al. (2001) for different line profiles for the northern and southern sources in 1202--0725. BRI1202--0725 is a radio continuum source with a total flux density of $315\\pm80$ $\\mu$Jy at 1.4 GHz (see section 4.1), and 141$\\pm$15 $\\mu$Jy at 4.9 GHz (Yun et al. 2000). Extrapolating the centimeter spectrum to millimeter wavelengths assuming a continuous powerlaw spectrum implies a non-thermal contribution to the integrated emission at 40 GHz of 80$\\mu$Jy. Conversely, extrapolating the (sub)mm dust emission spectrum downward to 40 GHz implies a thermal dust contribution of 25$\\mu$Jy, assuming $\\beta = 1.5$. Optical imaging of 1202--0725 shows a point-like QSO with M$_B$ = --28.5, with faint near-IR and Ly$\\alpha$ emission extending about 2.4$''$ north of the QSO (Hu, McMahon, \\& Egami 1996). The optical position of the QSO given by Hu et al. is 0.5$''$ north of the position of the southern mm/cm/CO component. In the following we assume that this offset indicates the accuracy of the relative astrometry of the radio and optical images. The (3$\\sigma$) limit to the Ly $\\alpha$ flux from the northern mm/cm/CO component in 1202--0725 is $\\sim 1\\times10^{-17}$ erg cm$^{-2}$ s$^{-1}$ (Hu et al. 196), corresponding to a limit to the Ly $\\alpha$ luminosity of $2\\times 10^{42}$ erg s$^{-1}$. For comparison, the Ly$\\alpha$ luminosities of UV dropout galaxies at $z \\sim 3.09$ are typically $\\rm 2~ to ~4 \\times 10^{42}$ erg s$^{-1}$ (Steidel et al. 2000). The line emitting object detected 2.4$''$ northwest of the QSO has a Ly$\\alpha$ luminosity of $4\\times 10^{43}$ erg s$^{-1}$ (Fontana et al. 1998; Petitjean et al. 1996). If this line emission is powered by star formation, then the star formation rate is $> 20$ M$_\\odot$ year$^{-1}$, depending on the amount of dust extinction (Fontana et al. 2000). It is also possible that this line emission is powered by UV radiation from the QSO itself (Petitjean et al. 1996). The lack of CO and thermal dust emission from this position argues for the latter, while the lack of N V and C IV emission lines argues for the former (Fontana et al. 1998; Ohta et al. 2000). \\subsection{1335--0417} Guilloteau et al. (1997) have detected CO (5--4) line emission from BRI~1335--0417 at $z = 4.4074$, with a line FWHM = $420$ km s$^{-1}$ and an integrated line intensity of 2.8$\\pm$0.3 Jy km s$^{-1}$. The flux density of the dust continuum emission measured at 1.35mm with the Plateau de Bure Interferometer (PdBI) at 2$''$ resolution is 5.6$\\pm$1.1 mJy, while that measured at 1.25mm with the IRAM 30m telescope (10.6$''$ resolution) is $10.3\\pm1.4$ mJy. The implied far IR luminosity based on the 1.25 mm measurement is $3.1 \\times 10^{13}$ L$_\\odot$. The 1.35mm PdBI observations shows marginal evidence for an extended source, with a formal Gaussian size of $1.0 \\pm 0.4''$, with major axis oriented roughly north-south. No high resolution optical images have been published of 1335--0417. The DSS shows an unresolved source with M$_B$ = --27.3, although with a pixel scale of just 1$''$ the details of the source structure remain unknown. BRI~1335$-$0417 has been detected in the radio continuum, with an integrated flux density of $220\\pm 43$ $\\mu$Jy at 1.4 GHz and $76\\pm 11$ $\\mu$Jy at 4.9 GHz (Carilli et al. 1999). ", + "conclusions": "We briefly compare the properties of 1202--0725 and 1335--0417 to those found in nuclear starburst galaxies seen at lower redshift. Yun et al. (2000) found that the ratio of radio-to-far IR luminosity for both these sources is within the range defined for active star forming galaxies based on the tight radio-to-far IR correlation (Condon 1992), although in both cases the ratio falls at the high end of the normal range, suggesting a possible contribution to the radio emission from the AGN. For 1202--0725 there is marginal evidence that the radio emission is time variable. Variability would rule-out a starburst origin for the radio emission. Further radio continuum monitoring is in progress to test this interesting possibility. The measured (redshift corrected) brightness temperature for the CO(2-1) emission from the compact components in the southern source in 1202--0725 is $\\ge 25$ K. This limit is consistent with values seen for the CO(2-1) emission from starburst nuclei of ULIRGs, for which brightness temperatures of 30 K to 60 K have been measured (Downes and Solomon 1998). The excitation conditions for the CO for these two sources follow roughly those seen for the nuclear starburst galaxy M82, but are very different than those expected for the disk of a normal spiral galaxy. Assuming a Galactic abundance for CO of $\\rm {{[CO]}\\over{[H_2]}} = 5\\times 10^{-5}$ (eg. Wilson et al. 1986) implies an H$_2$ column density of order 10$^{24}$ cm$^{-2}$. This value is similar to the molecular gas column density seen toward ULIRGs (Downes \\& Solomon 1998), but is an order of magnitude larger than that observed in the lower luminosity nuclear starburst galaxies M82 and NGC 253 (Harrison, Henkel, \\& Russell 1999, Mao et al. 2000). A low metallicity would suppress CO emission and would yield even higher H$_2$ column densities. It would be very interesting to measure the CO(1-0) line of the source. A direct comparison of CO(1-0) and 2-1 line intensities would provide significant information on optical depths, since the 1-0 line may be optically thin, while the 2-1 line is likely optically thick. We have found that the continuum-to-line ratio, ${{L_{\\rm FIR}}\\over{L'_{\\rm CO(1-0)}}}$, is about 335 for both sources. This ratio is at the high end of those seen for ULIRGs, and suggests a continuation of the non-linear trend for increasing continuum-to-line ratios with increasing far IR luminosity. Considering $L_{\\rm FIR}$ to be a measure of star formation rate, and $L'$(CO(1-0)) to be a measure of molecular gas mass, it has been suggested that this non-linear relation might imply a higher star formation efficiency ($\\equiv \\rm {{Star~Formation~Rate}\\over{Gas~Mass}}$) in higher luminosity galaxies, in particular for galaxies with $L_{\\rm FIR} \\ge 10^{11}$ L$_\\odot$ (Solomon et al. 1997). For dense nuclear starbursts a number of groups (Solomon et al. 1997; Mao et al. 2001; Weiss et al. 2001) have shown that the densities are such that the entire interstellar medium in the starburst regions may be molecular, and that the CO(1-0) emission may be dominated by this molecular inter-cloud medium, as opposed to being from the denser star forming clouds themselves. This phenomenon would contribute to the non-linear relationship between $L_{\\rm FIR}$ and $L'$(CO(1-0)). For 1202--0725 and 1335-0417 there is also the obvious possibility of dust heating by the AGN, in which case the continuum-to-line ratio cannot be interpreted in the context of star formation. Downes and Solomon (1998) show that for the nuclear starburst regions in ULIRGs the H$_2$ mass-to-CO(1-0) luminosity conversion factor, X, is a factor four or so below the Galactic disk value. Even for this low value of X, we find that the CO emitting regions in 1202--0725 and 1335--0417 must be close to face-on in order to avoid having the gas mass exceed the gravitational mass, implying perhaps unreasonably large rotational velocities. While this problem is mitigated somewhat by lowering X even further, the required X values become comparable to, or lower than, the minimum values dictated by optically thin CO emission (Solomon et al. 1997). One way of circumventing this mass problem would be to assume that the source is magnified by strong gravitational lensing. Magnification by a factor three or so would avoid unphysical values of X. The very large IR and CO luminosities, and the double structure of the sources in the CO line emission, and in the non-thermal radio continuum and thermal mm continuum for 1202--0725, could possibly indicate gravitational lensing. On the other hand, neither source is double at optical wavelengths, arguing against lensing, although the possibility of differential obscuration along the two lines-of-sight complicates this conclusion (Hu et al. 1996). More telling is the difference in CO(2-1) line profiles, which is difficult, although perhaps not impossible, to explain in the context of gravitational lensing. Also, double sources arise from strong gravitational lensing of very compact emitting regions (sizes $\\le$ few pc). For extended emitting regions ($\\ge 100$ pc), such as must be the case for the thermal CO and mm continuum emission, strong lensing will only occur if the extended emission regions cross a caustic in the source plane (Blandford \\& Narayan 1992). Such a phenomenon usually leads to more complex geometries, like arcs or rings, as in APM 0827+525 and IRAS 10214 (Lewis et al. 2001; Scoville et al. 1995). More sensitive, high resolution imaging at cm and mm wavelengths is required to address this interesting question. Overall, the physical conditions in the molecular gas and dust in these systems are similar to those observed in nuclear starbursts at low redshift, including: (i) the radio-to-far IR luminosity ratio, (ii) the CO brightness temperature, (iii) the CO excitation conditions, (iv) the CO column densities, and (v) the CO line-to-dust continuum ratio. \\vskip 0.2truein The National Radio Astronomy Observatory (NRAO) is operated by Associated Universities, Inc. under a cooperative agreement with the National Science Foundation. We thank E. Hu for allowing us to reproduce the Ly$\\alpha$ image, S. Myers for discussions concerning gravitational lensing, and the referee for many important comments." + }, + "0112/astro-ph0112347_arXiv.txt": { + "abstract": "{This paper presents an automated method to determine detailed abundances for A and F-type stars. This method is applied on spectra taken with the ELODIE spectrograph. Since the standard reduction procedure of ELODIE is optimized to obtain accurate radial velocities but not abundances, we present a more appropriate reduction procedure based on IRAF. We describe an improvement of the method of Hill \\& Landstreet (\\cite{hill}) for obtaining $V\\sin{i}$, microturbulence and abundances by fitting a synthetic spectrum to the observed one. In particular, the method of minimization is presented and tested with \\object{Vega} and the \\object{Sun}. We show that it is possible, in the case of the \\object{Sun}, to recover the abundances of 27 elements well within 0.1 dex of the commonly accepted values. ", + "introduction": "The determination of detailed abundances requires a high resolving power ($> 30\\,000$) and a wide spectral range. In order to satisfy both requirements simultaneously, echelle spectrographs must be used. ELODIE (Baranne et al. \\cite{baranne}, hereafter BQ96) is a fiber-fed echelle spectrograph with a resolution of R=42\\,000 attached to the 1.93\\,m telescope of the Observatoire de Haute-Provence (OHP), France. This spectrograph and its reduction software were optimized to measure accurate radial velocities. In this paper we first show what precautions have to be taken to use ELODIE for other spectroscopic analyses, in our case detailed abundance determinations. To achieve our goal we had to make another reduction, starting from the raw image and taking special care in the removal of scattered light. Another important point in the reduction is to paste together the different orders of the spectrum and normalize them. Secondly, we present a method to estimate abundances with synthetic spectra adjustments. This method is an improvement of that of Hill \\& Landstreet (\\cite{hill}, HL93 hereafter). It is automated as much as possible and is able to analyse stars with various rotational velocities (up to 150 $\\mathrm{km~s^{-1}}$), for which the equivalent width method is not applicable. Finally, to assess the validity of this method, we compare the abundances derived for \\object{Vega} (\\object{$\\alpha$ Lyr} = \\object{HR 7001} = \\object{HD 172167}) and the \\object{Sun} with those in the literature. These two reference stars are used to check the code's validity for stars having effective temperatures between those of the \\object{Sun} and of \\object{Vega}. Analysis tools with related goals but different approaches have been developed by Valenti and Piskunov (\\cite{valenti}), Cowley (\\cite{cowley}) and Takeda (\\cite{takeda}). Takeda's method has been used by Varenne and Monier (\\cite{varenne}) to derive abundances of A and F-type stars in the Hyades open cluster. ", + "conclusions": "We have shown that our method of determining detailed abundances for A-F type stars works and that the ELODIE echelle spectrograph can be used to get accurate abundance determinations for \\object{Vega} and the \\object{Sun}. We achieve a high level of automation to extract spectra and analyze them. One line list covering the ELODIE spectral range was compiled and will be used for the following study of A-F stars. It includes new gf values for some critical lines determined using the solar spectrum. The problem of the choice of stellar parameters which arose during the analysis of Vega is also a good justification to analyze a large sample of stars with one given method for determining stellar parameters ($T_\\mathrm{eff}, \\log{g})$. The choice between a photometric or spectroscopic method is not so important since the uncertainty of these methods is comparable. Next it is important to determine abundances of all stars of the sample in a homogeneous way, and that will be possible with the method presented in this paper. This is our final goal for which automation will be crucial. Therefore, even if some uncertainties remain, the resulting errors will be systematic, and will not depend on the author's subjectivity. Finally, our work is further justified by the commissioning of medium and high resolution multi-fibers spectrographs, because when an observer gets hundreds of spectra each night, he can no longer handle them by hand." + }, + "0112/astro-ph0112037_arXiv.txt": { + "abstract": "We have observed the young Galactic supernova remnant Kes~75 with the Chandra X-ray Observatory. This object is one of an increasing number of examples of a shell-type remnant with a central extended radio core harboring a pulsar. Here we present a preliminary spatially resolved spectroscopic analysis of the Kes~75 system. We find that the spectrum of the pulsar is significantly harder than that of the wind nebula, and both of these components can be isolated from the diffuse thermal emission that seems to follow the same distribution as the extended radio shell. When we characterize the thermal emission with a model of an under-ionized plasma and non-solar elemental abundances, we require a significant diffuse high energy component, which we model as a power-law with a photon index similar to that of the synchrotron nebula. ", + "introduction": "Kes 75 (also known as G29.7$-$0.3) is one example in our Galaxy of a young, shell-type remnant ($3.5\\arcmin$ in diameter) with a central core ($30\\arcsec$) whose observed properties suggest a synchrotron nebula similar to the Crab Nebula (Becker, Helfand \\& Szymkowiak 1983; Becker \\& Helfand 1984; Blanton \\& Helfand 1996). Observations by the Advanced Satellite for Cosmology and Astrophysics (\\asca) verified the existence of both thermal and non-thermal emission, but lacked the spatial resolution necessary to separate the components (Blanton \\& Helfand 1996). Monitoring with the Rossi X-Ray Timing Explorer led Gotthelf et al. (2000) to discover a 700 year old pulsar, PSR~J1846$-$0258, located in the \\asca\\ data within the Crab-like core. Expanding on the previous work, we discover a significant spectral difference between the pulsar and the surrounding wind nebula. Using a non-equilibrium ionization model we find different temperature and heavy-metal abundances in separate regions of the shell, as well as evidence of non-thermal emission throughout. ", + "conclusions": "For the \\asca\\ observation of Kes~75, it was not possible to extract only the spectrum of the central Crab-like component of the supernova remnant without the surrounding thermal part, only to fit a model to all the data that accounts for both of the components. Given this limitation, Blanton \\& Helfand (1996) characterized all the non-thermal emission of the object with a single power-law ($\\Gamma = 2.0$). We can now resolve the core from the shell, and analyze their spectra individually. The clearly non-thermal spectrum of the central nebula exhibits a photon index consistent with the \\asca\\ findings, yet if this were the sole source of non-thermal flux, our heuristic model of the shell portion should be just as consistent with the previous results. The lack of similarity between the two spectra indicates a significant amount of non-thermal flux from outside the central core of the remnant. This extra emission may account for the spectral inconsistencies between the line strengths and the thermal continuum." + }, + "0112/gr-qc0112070_arXiv.txt": { + "abstract": "\\noindent What is the nature - continuous or discrete - of matter and of its fundamental interactions? The physical meaning, the properties and the consequences of a discrete scalar field are discussed; limits for the validity of a mathematical description of fundamental physics in terms of continuum fields are a natural outcome of discrete fields with discrete interactions. Two demarcating points (a near and a far) define a domain where no difference between the discrete and the standard continuum field formalisms can be experimentally detected. Discrepancies, however, can be observed as a continuous-interaction is always stronger below the near point and weaker above the far point than a discrete one. The connections between the discrete scalar field and gravity from general relativity are discussed. Whereas vacuum solutions of general relativity can be retrieved from discrete scalar field solutions, this cannot be extended to solutions in presence of massive sources as they require a true tensor metric field. Contact is made, on passing, with the problem of dark matter and the rotation curve of galaxies, with inflation and the accelerated expansion, the apparent anomaly in the Pioneer spacecraft acceleration, the quantum Hall effect, high-$T_{c}$ superconductivity, quark confinement, and with Tsallis generalized one-parameter statistics as some possible manifestation of discrete interaction and of an essentially discrete world. ", + "introduction": "Although a scalar field has not been observed in nature as a fundamental field its use as such is very frequent in the modern literature, particularly in elementary particles, field theory and cosmology. Here we will apply to the scalar field the concepts and results developed in the reference \\cite{hep-th/0006237}, where the concept of a discrete field was introduced and its wave equation and its Green's function discussed. The standard field and its formalism, which for a distinction, we always append the qualification continuum, are retrieved from an integration over the discrete-field parameters. Remarkable in the discrete field is that it has none of the problems that plague the continuum one so that the meaning and origin of these problems can be left exposed on the passage from the discrete to the continuum formalism. Although the motivations for the introduction of a generic discrete field in \\cite{hep-th/0006237} have being made on pure physical grounds of causality, a deeper discussion about its physical interpretation have been left for subsequent papers on specific fields. This discussion will be retaken here with the simplest structure of a field, the scalar one. The idea of a pointlike field, although unusual, represents the same symmetry of quantum field theory where fields and sources are equally treated as quantized fields. Here it is seen from a classical perspective. Besides, pointlike object is not a novelty in physics and one of the major motivations of the, nowadays so popular, string theory is of avoiding \\cite{Polchinski} infinities and acausalities in the fields produced by point sources, problems that do not exist for the discrete field \\cite{hep-th/9610028}. This paper is structured in the following way. Section II, on the sake of a brief review of the mathematical definition of discrete fields, is a recipe on how to pass from a continuum to a discrete field formalism, and vice-versa. The discrete scalar field, its wave equation, Lagrangian and energy tensor are briefly discussed in Section III. Section IV discusses the consequences of discrete interactions for the mathematical description of the physical world. Then it gains generality as the discussions leaves the specificity of scalar interactions widening to the universality of all fundamental interactions. Calculus (integration and differentiation) which is based on the opposite idea of smoothness and continuity, has its full validity for describing dynamics restricted then to a very efficient approximation in the case of a high density of interaction points, such that the concept of acceleration as a continuous change of velocity may be introduced in an effective physical description of fundamental interactions. This seems to be an answer to the Wigner's pondering \\cite{Wigner} about the reasons behind the unexpected effectiveness of mathematics on the physical description of the world. In Section V the discrete scalar field reveals and unavoidable connection to gravity of general relativity as both have energy-momentum for sources. Besides, it has been shown \\cite{gr-qc/9801040} that the gravitational field of general relativity in a vacuum is reduced to an effective spacetime manifestation of a discrete scalar field. We will see here that this cannot be extended to gravity in the presence of massive sources; a true second-order tensor field is required. The reader should be aware, however, that this is not a paper about a new theory of gravity; it is about the meaning and implications of an assumed discreteness of all matter and interactions, particularly of the scalar ones. Its connection to gravity is a consequence. Theoretical implications of discrete interactions and the possibility that some of its consequences may have already been observed are considered in the concluding Section VI. ", + "conclusions": "The thesis that all fundamental interactions are discrete is being developed. If this is the case there is no really compelling reason for excluding gravity from such a unifying idea. The knowledge of a supposedly true discrete character of all fundamental interactions is a permanent reminder of the limits of a continuum approximative description. There are many indications of possible theoretical and observational evidences of direct consequences of interaction discreteness, particularly in gravity, of which we can just make some brief comments, although they are full-attention deserving. Discrete interactions imply on the ratio parameter (\\ref{ratio}) whose range is divided in three segments by its two (the far and the near) critical values. The interior segment, between these two values, defines the domain of validity of the continuous-interaction approximation where the polygonal worldline of the sources are so densely packed of interaction points that they can be effectively replaced by smoothly continuous curves where the concept of acceleration and of spacetime curvature at a point on the worldline make sense. On the other hand, for values of the ratio parameter (\\ref{ratio}) in the exterior segments, i.e. below or above the two critical points the discrepancies between a discrete and a continuous interaction cannot be overlooked. These are the places where great riddles are associated to gravity. They are not easily fitted with a continuous interaction without evoking new forms of interaction or of (dark) matter and energy. This casts doubts on the results about asymptotical fields and their singularities of any continuum-field theory. \\subsection{Large distance effects} Gravity, in contradistinction to all other fundamental interactions has problems in the large distance limits where the field is extremely weak. The continuum asymptotically null fields are replaced by discrete interactions that become more and more sparse with the distance. This may be detectable for the gravitational field as it does not have shielding effects although it requires huge masses for detecting very weak gravitational fields and huge distances for producing a detectable $\\Delta\\tau_{j}$; both conditions found at and above galactic scales. Therefore, a right place for checking for signs and applications of discreteness is on dynamics of galaxies and on cosmology. Flat rotation curve of galaxies, inflation, accelerated expansion may be just manifestation of a discrete gravity. Asymptotic flat rotation curves appear very naturally \\cite{gr-qc/0103218} in a discrete-field context! It is therefore a real possibility that the critical point for gravity has already been detected in the flat rotation curves of galaxies \\cite{rotation curves of galaxies}. The flatness feature of a rotation curve of a galaxy, as remarked by Milgrom \\cite{Milgrom}, is determined not by its central mass $M$ alone nor just by the distance $R$ but by the acceleration which is equivalent to the ratio-parameter (\\ref{ratio}) as $\\Delta q_{j}$ corresponds to a change of speed. Therefore the existence of the critical point in the continuum/discrete physical description justify the introduction of a new fundamental scale for effective accelerations, and may put Milgrom's MOND on a sound physical basis. The actually prevailing wisdom that a flat rotation curves is the (ad hoc) indication of some strange, ubiquitous but still to be detected cold dark matter is not free of problems and is far from being unanimous\\cite{Milgrom,cdm,Evans,Mannheim,Nucamendi}. A gravitational repulsion, instead of attraction may be just a residual correction of the excess committed in the continuum approximation. In a scale of distance where the effective attractive interaction is so weak that sub-dominant contributions to discrete gravity dominate \\cite{gr-qc/0111152}. Another possible evidence of discrepancy that must be considered is the apparent anomalous, weak, long-range acceleration observed in the Pioneer 10/11, Galileu, and Ulysses data \\cite{Pioneer}. Due to their spin-stabilization and to the great distance (30 t0 67 AU) from the Sun the spacecraft are excellent for dynamical astronomy studies as they permit precise acceleration estimation to the level of $10^{-10}cm/s^{2}.$ The detected anomalous acceleration comes from neglected second largest contribution from those mentioned $n$-combinatorials, although it is still not clear if this anomaly is not a too large contributions from discrete interactions. \\subsection{Small distance effects} In the asymptotical short-distance limit, below the critical point, the fact that the discrete field is weaker than the continuum one suggests a natural explanation to inflation in cosmological theories. Other possible indication of discrete physics are the experimentally observed (or just realized) manifestations of low (one- or two-) dimension physics. It is an enormous difficulty for the continuum field formalism to explain a mechanism for restricting the interaction field to a one- or two-dimensional space instead of spreading over the whole three-dimensions. Quantum Chromodynamics faces similar problems for explaining quark-confinement. Solid state physics is, in a discrete interaction context, a totally unexplored subject, rich of such possible manifestation of discrete physics. Possible examples are the quantum Hall Effect and the high $T_{c}$ superconductivity where indications of two-dimensional physics have been observed. Crystal or molecular structural arrangements constraining the movement of the relevant charges to layers (subdimensional regions) restrict, as a consequence, the discrete, but not the continuum, interaction fields to these sub-regions. Whereas this is easily explained with discrete fields it does not work for the continuum one as it would require ultrarelativistic velocities for an approximate (but in the wrong space direction) dimensional reduction of the field 3-dimensional support. This same mechanism explains, in terms of discrete fields, light polarization and the action of a polaroid. These constraints may be dependent on the conditions of low temperatures with its consequent suppression of some thermal excitations. Since there is no differential equations and neither integrations with discrete interactions, the evolution of any system is done through successive, sudden and discrete finite differences that are just super imposed. Between two consecutive interaction points every pointlike component just moves freely on straight lines. It is then remarkable that the whole continuum-field physics is reproduced as an effective approximation after a large number of discrete interactions \\cite{gr-qc/0111152}. Any exact physical statement must be expressed as finite power series of combinatorials of the accumulated number of interaction events. So the world is surprisingly simpler and our standard vision of it is richer of such idealized, unreachable concepts than we had previously conceded. A whole paraphernalia of mathematical tools, so useful in physics - differential equations, integrations, differential geometry, topology, just for citing few - and also so many familiar and daily used mathematical functions like sine, exponential, harmonic and coulombian potentials, and circles, ellipse, etc, do not belong to the realm of the physical world; they are just unreachable limiting concepts as much as an ideal gas and a reversible process. This is reflected in the Tsallis's \\cite{Tsallis,Tsallis's homepage} generalized one-parameter statistics based on a power-law distribution of probabilities. It is reduced to Boltzmann statistics when its parameter is equal to one. This parameter is then a measure of how close the system is from its idealized asymptotic state, that rigorously, is reachable only after an infinite number of interactions. It may be an appropriate statistics for a world made of sets of discretely interacting pointlike objects. \\subsection{Concluding remarks} The idea of an essential continuity of any physical interaction allows unlimited speculations that will always go beyond any level of possible experimental verifications which brings then the risk of not being able of distinguishing the reign of possibly experimentally-grounded scientific research from plain philosophical speculation or even just fiction. Regardless the possibility that some of its consequences have already been experimentally detected, a discrete gravitational interaction, even in other range where it is not experimentally detectable, still for a long time to come, may just make sense of existing theories for delimiting their domain of validity as it has historically happened with all new discreteness introduced in the past, like the ideas of molecules, atomic transitions, and quarks, for example." + }, + "0112/astro-ph0112201_arXiv.txt": { + "abstract": "{ The TeV emission of low power BL Lac objects has been established by the detection of an handful of them. The knowledge of the level of the TeV emission and its spectrum can shed light on the particle acceleration mechanisms, and it is especially important to assess the still uncertain level of the far infrared background radiation, which can absorb the TeV photons through photon--photon interactions. In view of these implications, it is necessary to enlarge the number of TeV detected sources, and to find them at different redshifts. To this aim, we propose a general and simple criterium to select the best TeV candidates, and produce a list of them with flux estimates above 40 GeV, 300 GeV and 1 TeV. ", + "introduction": "The discovery that blazars (Flat Spectrum Radio Quasars and BL Lac objects) are very strong $\\gamma$--ray emitters have renewed the interest about them, and opened new perspectives in the comprehension of the physics of these objects. The observations by EGRET (Hartmann et al. 1999), onboard the {\\it Compton Gamma Ray Observatory}, led to the discovery that blazars emit most of their power in the $\\gamma$--ray band, and that their Spectral Energy Distribution (SED) is characterized by two broad peaks, now commonly (but not unanimously), interpreted as due to synchrotron and inverse Compton (IC) radiation, respectively (e.g. Maraschi et al. 1992; Dermer et al. 1992; Sikora et al. 1994; Ghisellini \\& Madau 1996; but see Mannheim 1993; Rachen 1999; Muecke \\& Protheroe 2000; Aharonian 2000 for a different interpretation). For the first time we were able to study their entire SED in a comprehensive way, finding differences among subclasses of blazars about the frequency location of the two broad peaks, their relative luminosity and variability behaviors in different bands. Considering the EGRET sources and three complete blazar samples, it was found a correlation between the location of the two broad peaks and the observed bolometric luminosity (Fossati et al. 1998, hereafter F98; Ghisellini et al. 1998). Blazars seem to form a well defined sequence, with low powerful objects having both peaks at a similar level of luminosity, and located at higher frequencies than in more powerful objects, in which the IC peak dominates the emission. In the BL Lac class, the first kind of sources were named High frequency Peak BL Lacs (or HBL for short) by Padovani \\& Giommi (1995), while the latter subclass was called Low frequency Peaked BL Lacs (LBL). Recent observations of high redshift ($z>4$) blazars (Fabian et al. 2001a, 2001b; Celotti 2001), and of low power BL Lac objects (Costamante et al. 2001a, 2001b) have extended the blazar sequence at both ends, resulting in agreement with the original trend. At TeV energies, the detection and study of blazar objects by ground based Cherenkov telescopes have been limited up to now to few HBL sources (Mkn 501, Mkn 421, PKS 2155--304, 1ES 2344+514, see Catanese \\& Weekes 1999), though disclosing new and fundamental aspects of the blazar behavior. These observations monitor the behavior of the most energetic electrons of the source, thus shedding light on the acceleration mechanism working at the most extreme conditions. The very rapid variability observed at these energies (Mkn 421 doubled its TeV flux in less than 20 minutes, see Gaidos et al. 1996, Catanese \\& Weekes 1999), coupled with the requirement that the source must be transparent with respect to the photon--photon process, tightly constrains the physical parameters, such as the source size and its beaming factor. In addition, Mkn 421 showed a tight correlation between the emission in the X--ray and TeV bands (Maraschi et al. 1999, Takahashi et al. 2000, Krawczynski et al. 2001), implying that the radiation produced in the two bands is co-spatial and produced by the same electrons: this is of crucial importance to constrain any emission model. The strong connection between the TeV and X-ray emission was also clearly evident during the 1997 flare of Mkn 501, when this source was observed by the X-ray satellite {\\it Beppo}SAX in an extreme spectral state, with a synchrotron peak frequency close to 100 keV or even more (Pian et al. 1998). Mkn 501 was found to have increased at least tenfold its luminosity, with most of it radiated at high X--ray energies. At the same time, the source underwent a major flare in the TeV band (Catanese et al. 1997b; Aharonian et al. 1997; Protheroe et al. 1997; Djannati--Ata\\\"{\\i} et al. 1999), and continued to be active (and well visible) in the TeV band for several months. This dramatic behavior can be explained by a synchrotron inverse Compton model, taking into account the effects introduced by the Klein Nishina scattering cross section and the constraints posed by the transparency of the source with respect to photon--photon collisions producing electron--positron pairs. It is not clear if the simultaneous variations in the X--ray and TeV bands can be completely accounted for by a simple one--zone homogeneous synchrotron self--Compton model (see Tavecchio et al. 2001), or if we need some extra and more quiescent source of IR--optical photons (i.e. Ghisellini 1999). TeV observations of BL Lacs are also particularly interesting because, being the only extragalactic sources known to emit at these energies, allow an independent estimate of the extragalactic IR background (IRB). Direct measurements of the IRB are affected (up to now) by relatively large uncertainties, due to the heavy contamination of foreground objects (for a review see Hauser \\& Dwek 2001). Because TeV photons can be absorbed by IR photons for the pair production mechanism, the analysis/study of high energy spectra from sources at different redshifts gives an independent measure of the IRB level (Stecker et al. 1992). These kind of studies are just started, and the first conclusions are based on the two most observed TeV BL Lacs, namely Mkn 421 and Mkn 501. The main uncertainty here is the knowledge of the primary blazar spectrum, which could have an intrinsic cut--off at high energies, or could be affected by absorption due to IR photons produced locally. A first and preliminary confirmation of the IRB absorption has been obtained comparing the spectra of Mkn 421 and Mkn 501, showing a cut--off in the power--law spectra at approximately the same energies (Krennrich et al. 2001), as expected since they have similar redshifts. The direct measurements of the IRB flux lead to predict quite a strong absorption at TeV energies. If true, this in turn would imply an unusual primary spectrum above $\\sim 1$ TeV, which must have an excess above the extrapolation from lower energies, leading to the so--called IRB--TeV puzzle (see e.g. Protheroe \\& Meyer 2000; Aharonian, Timokhin \\& Plyasheshnikov, 2001; Berezinsky 2001) {\\footnote{It has been proposed that this ``puzzle\" could be solved by quantum--gravity theories predicting the breaking of Lorentz invariance: as a consequence there should be a modification in the energy threshold for the $\\gamma$--$\\gamma$ $\\to$ $e^\\pm$ process, explaining why TeV photons are not heavily absorbed by the IRB (Amelino--Camelia \\& Piran 2001).}. However, to draw unambiguous conclusions, we need an ensemble of sources located at different redshifts and a detailed knowledge of the X--ray flux and spectrum: being produced by the same electrons (in synchrotron inverse Compton models), this would help in predicting the shape of the TeV emission. To this aim we consider in this paper several published samples of bright BL Lac objects, for a total of 246 different objects, and propose a simple and handy tool to identify and select the most promising candidates. The main point we emphasize concerns the requirement of {\\it both} high energy electrons {\\it and} sufficient seed photons to originate the TeV emission. We therefore consider as best candidates those BL Lac objects having not only their synchrotron peak located at high energies, but also having sufficient radio--through--optical flux. We therefore expand the work first done by Stecker et al. (1996), concerning only the Einstein Slew survey sample of BL Lac objects, both by considering other samples and also by introducing a different, albeit still simple, criterium to identify the best candidates. \\begin{figure*} \\psfig{figure=sed_tev.ps,width=18cm,height=18cm} \\vskip -0.8 true cm \\caption{Spectral energy distributions of the 5 BL Lac objects already detected at TeV energies. The solid and dashed lines refer to the SSC model described in Section 3.2 and to the SED constructed using the parameterization described in Fossati et al. 1998, with the modification described in this paper. The input parameters used for the SSC models are given in Ghisellini, Celotti \\& Costamante (2001). } \\label{tevdet} \\end{figure*} ", + "conclusions": "With new Cherenkov telescopes foreseen to operate in the next few years the TeV extragalactic astronomy is entering its adulthood. A tenfold increase in sensitivity (expected for the forthcoming installations) would mean the possible detection of $\\sim$ 100 BL Lacs, if the counts of BL Lac object at TeV energies are roughly Euclidean (in the bright flux end) and neglecting absorption by IRB. Because of IRB absorption, the counts will be flatter than Euclidean, but the lower energy threshold of some new instruments may compensate % for the extragalactic absorption, as well as favoring the detection of slightly less blue objects. Therefore many more sources are expected to be detectable by the new telescopes, and this motivated us to study which kind of BL Lac objects is more likely to be detected at high energies. Our findings can be summarized very simply once we realize that, to produce a strong TeV emission, the inverse Compton process needs a sufficient number of both very high energy electrons and soft seed photons. Therefore we require both a strong X--ray flux and a sufficiently strong radio--through--optical flux. Since the optical flux can be contaminated, especially in low redshift sources, by the underlying host galaxy, our sources are primarily selected as bright both in the X--ray and radio bands. All but one of these sources are also bright in the optical. With respect to the previous work by Stecker et al. (1996), our criterium introduces the further requirement that the source must be a relatively strong radio emitter. The other difference is that we considered not only the Einstein Slew survey sample of BL Lacs, but several other BL Lac samples. Besides selecting the best candidates through their location in the radio -- X-ray flux plane (a criterium largely model independent), we have also tried to quantify the level of the expected high energy emission for each selected source, by applying a one--zone synchrotron self--Compton model and also the phenomenological description of the SED of Fossati et al. (1998), slightly modified to better account of the average SED of low power BL Lacs. The latter model, by construction, assumes equal power between the synchrotron and the inverse Compton components of the SEDs, and almost always predicts larger high energy fluxes than the SSC model. In the SSC model, in fact, the Compton dominance is not fixed a priori, but found by fitting the synchrotron part of the spectrum, which fixes the value of the magnetic field. We stress that the F98 prescription was designed including also the SEDs of those BL Lacs already detected in the TeV band (and indeed is in good agreement with their TeV flux levels), and therefore seems more appropriate to predict the TeV flux of sources in high state. The adopted SSC model, instead, is designed to fit the known synchrotron part of the SED, which is often representative of a more ``normal\" or quiescent state. This explains the sometimes large discrepancy between the predicted fluxes of the two models. Since BL Lac objects are among the most variable sources, especially at high energies, the two foreseen flux levels could be thought of as an approximate range of variability, and the average flux could be considered as a measure of the probability to find the source in a particular TeV state. We would like to stress, anyway, that the uncertainties on the key parameters we used for the model and the non--simultaneity of the fitted data lead to large uncertainties in the predicted TeV flux, sometimes of the same order of the differences between the two adopted models. Our predicted fluxes, therefore, also in the case of the SSC model, must be considered as ``best guesses\" on the high energy emission from these objects. Within these limits, the SSC model provides more information than the phenomenological parameterization, since it gives also the expected shape of the high energy spectrum. Since the level of the synchrotron X--ray flux measures, in our scenario (as well as in the Stecker et al. 1996 one), the number of TeV energy electrons, the X--ray monitoring of our candidates is particularly useful to catch sources in high TeV states, as already partially done through the All Sky Monitor (ASM) onboard the {\\it Rossi}XTE satellite. Besides the fluxes above 300 GeV, the most common threshold of present Cherenkov telescopes, we also give our estimates above 40 GeV, which is approximately the energy threshold of CELESTE and of forthcoming observatories like HESS and MAGIC (with VERITAS at $\\sim50$ GeV, Weekes 1999). Emission at these energies is much less absorbed by the cosmic infrared background, giving the opportunity to see more distant sources and study their intrinsic spectrum in an unabsorbed band. In addition this energy range will link ground based Cherenkov observations and the data coming from satellites, such as AGILE and GLAST, observing from a few tens of MeV to a few tens of GeV. As a final note, we warn that our flux estimates \\emph{do not} include the possible absorption due to the infrared background, since we preferred to be independent of this factor. In fact, we focussed on the conditions for the TeV emission, providing a list of possible sources, in order to allow an independent test of the IR absorption effects. This differs with respect to the flux estimates in Stecker et al. (1996), which instead account for the IR background absorption. However, since most of the photon flux is expected to be emitted at energies below 1 TeV (as can be seen in Table 3, comparing the values above 0.3 and 1 TeV), the reported fluxes should not be much affected for sources up to $z\\sim 0.1$, according to the present estimates on the IR background (see Stecker 2001 and references therein)." + }, + "0112/astro-ph0112567_arXiv.txt": { + "abstract": "\\xmm\\ was used to observe two eclipsing, magnetic cataclysmic variables, DP Leo and WW Hor, continuously for three orbital cycles each. Both systems were in an intermediate state of accretion. For WW Hor we also obtained optical light curves with the \\xmm\\ Optical Monitor and from ground-based observations. Our analysis of the X-ray and optical light curves allows us to constrain physical and geometrical parameters of the accretion regions and derive orbital parameters and eclipse ephemerides of the systems. For WW Hor we directly measure horizontal and vertical temperature variations in the accretion column. From comparisons with previous observations we find that changes in the accretion spot longitude are correlated with the accretion rate. For DP Leo the shape of the hard X-ray light curve is not as expected for optically thin emission, showing the importance of optical depth effects in the post-shock region. We find that the spin period of the white dwarf is slightly shorter than the orbital period and that the orbital period is decreasing faster than expected for energy loss by gravitational radiation alone. ", + "introduction": "\\begin{table*} \\caption[Observations summary]{ Details of the observations (see also Section \\ref{obsred}) } \\begin{tabular}{llcrrccccc} \\hline Object & Instrument & \\multicolumn{2}{c}{Beginning of} & Duration & Filter & Instrument & Timing & Aperture & Event \\\\ & & \\multicolumn{2}{c}{Observation (UTC)} & & & mode & resolution & radius & pattern \\\\ \\hline DP Leo & EPIC MOS 1 & 22 Nov 2000 & 4:55:02 & 22340 s & Thin 1 & full window & 2.6 s & 25'' & 0--12 \\\\ & EPIC MOS 2 & '' & 4:54:54 & 22348 s & Thin 1 & full window & 2.6 s & 25'' & 0--12 \\\\ & EPIC PN & '' & 5:36:22 & 20034 s & Thin 1 & full window & 0.073 s & 18'' & 0--12 \\\\ \\hline WW Hor & EPIC MOS 1 & 4 Dec 2000 & 3:37:25 & 23543 s & Thin 1 & full window & 2.6 s & 25'' & 0--12 \\\\ & EPIC MOS 2 & '' & 3:37:20 & 23548 s & Thin 1 & full window & 2.6 s & 25'' & 0--12 \\\\ & EPIC PN & '' & 4:18:48 & 21149 s & Thin 1 & full window & 0.073 s & 18'' & 0--12 \\\\ & Optical Monitor & '' & 3:30:09 & 9 $\\times$ 2200 s & B & fast timing & 0.5 s & 3.2'' & --- \\\\ & SAAO (1.0 m) & 2 Dec 2000 & 19:41:35 & 7140 s & R & --- & 60 s & --- & --- \\\\ \\hline \\end{tabular} \\label{obstable} \\end{table*} DP Leo and WW Hor are two eclipsing binary systems of type AM Her, a class of magnetic cataclysmic variables also referred to as polars because of their strongly polarized optical emission. In these systems the strong magnetic field of the white dwarf primary causes it to rotate synchronously with the orbital motion of the binary. Due to the strong magnetic field, the accretion stream from the Roche lobe filling secondary does not form a disk, but rather follows the magnetic field lines onto the white dwarf's surface. Slightly above the photosphere the accretion flow forms a shock that heats the gas to temperatures above $\\sim$10$^8$K. The gas below the shock front then cools and settles onto the photosphere \\citep{2000SSRv...93..611W}. Models of the accretion region suggest three major spectral components. Thermal bremsstrahlung is emitted in the X-ray band by the hot gas between the shock front and the surface. The photosphere below the shock is heated to temperatures of $\\sim$10$^5$K by reprocessing of bremsstrahlung photons and by dense filaments in the accretion stream \\citep{1982A&A...114L...4K}, giving rise to a blackbody component that extends from the UV to the soft X-ray band. The strong magnetic field at the surface of the white dwarf causes the shock-heated electrons to emit strongly polarized cyclotron radiation in the optical and IR bands. A comprehensive review of polars is given in \\citet{1990SSRv...54..195C}. Much has been learned about the accretion region in polars from optical polarimetry. These observations, however, are only able to reveal the physical conditions that give rise to the cyclotron radiation. To fully understand the structure of the post-shock region and to test dynamical models of the accretion flow, the thermal bremsstrahlung and the blackbody component, both visible in the X-ray band, need to be studied as well. Past X-ray observations were limited by low sensitivity and a narrow X-ray bandpass. With the high sensitivity of \\xmm\\ and its coverage of the 0.1--12 keV band we can now obtain light curves and spectra with unprecedented quality. This will allow us to independently study the bremsstrahlung and the blackbody component and derive the physical and geometrical properties of the accretion regions. In this paper we analyze X-ray and optical light curves of two polars in order to constrain the geometry of the accretion regions, investigate properties of the X-ray emitting gas, determine parameters of the binary systems, and find changes in the accretion longitudes and orbital periods. We present a detailed spectral analysis of these data in \\citet{2001MNRAS.326L..27R}. DP Leo, originally named E1114+182, was the first known eclipsing polar. This binary system, which has an orbital period of 89.8 min, was discovered serendipitously with the \\einstein\\ observatory and quickly identified as a polar because of its strongly modulated, polarized emission \\citep{1985ApJ...293..303B}. \\citet{1994ApJ...437..436R} and \\citet{1993MNRAS.261L..31B} later found that the accretion spot longitude is changing by $\\sim$2$^\\circ$ per year, suggesting a slightly asynchronous rotation of the white dwarf. We have shown in \\citet{2001MNRAS.326L..27R} that the X-ray spectrum obtained with \\xmm\\ contains in addition to the soft blackbody component a previously undetected hard bremsstrahlung component. Using the X-ray spectrum we derived a white dwarf mass of $\\sim$1.4$\\ $M$_\\odot$. WW Hor (EXO 023432-5232), which has an orbital period of 115.5 min, was discovered as a serendipitous X-ray source with \\exosat\\ and later identified as a polar \\citep{1987A&A...175L...9B}. Some of the system parameters were derived by \\citet{1988MNRAS.234P..19B}, but orbital inclination and accretion spot latitude could not be determined. A significant change of the accretion longitude over time was found by \\citet{1993MNRAS.261L..31B}. As we have shown in \\citet{2001MNRAS.326L..27R}, the \\xmm\\ spectrum does not contain the blackbody component typically found in polars at soft X-ray energies. From the X-ray spectrum we derived a white dwarf mass of $\\sim$1.1$\\ $M$_\\odot$. ", + "conclusions": "We have performed extensive \\xmm\\ observations of two magnetic cataclysmic variable stars, DP Leo and WW Hor. From a comparison with previous observations we conclude that both objects were in an intermediate state of accretion. We derived for both systems the mass ratio, orbital inclination and an improved eclipse ephemeris. We examined the X-ray data for periodic and quasi-periodic oscillations, but found no evidence for them in either source. \\subsection{DP Leo} From the shape of the hard X-ray light curves we conclude that the post-shock region has a significant optical depth. Although the X-ray spectrum indicates optically thin bremsstrahlung, it might still be consistent with an optically thick post-shock region if electron scattering is the major source of opacity. We find a strong asymmetry in the shape of the soft X-ray light curve which is not seen for the hard X-rays. This asymmetry can be explained with photoelectric absorption by the accretion curtain ($N_H\\sim 10^{20}{\\rm cm}^{-2}$) or with two additional accretion regions near the lower rotational pole. By comparing our measurement of the accretion spot longitude with previous results, we find that the longitude is changing linearly in time by $2.3^\\circ$ per year. This is likely due to the rotational period of the white dwarf being shorter than the orbital period by $\\sim$1$\\times10^{-6}\\ $s$\\ $s$^{-1}$. From a comparison with previous results we find that our measurement of the time of eclipse is not consistent with a linear ephemeris and that instead the orbital period of the system is decreasing on a time scale $P_{orb}/\\dot{P}_{orb}=-3.3\\cdot10^7\\ $yrs. This change is one order of magnitude faster than expected for energy loss by gravitational radiation alone. \\subsection{WW Hor} We use the shape of the X-ray light curves to derive the horizontal and vertical extent of the post-shock region. We find that the soft X-rays originate from a region that has a longitudinal extent of $\\sim$40$^\\circ$ while the hard X-rays are emitted in a region about half this size. This result shows that the temperature of the post-shock gas is increasing toward the centre. We find that the cyclotron radiation originates from a height of $\\sim$1\\% of the white dwarf radius while the X-rays are emitted at a much lower height. This is consistent with the general prediction that the temperature decreases and the density increases as the gas settles onto the surface. By comparing our measurement with previous results we find that the accretion longitude is correlated with the accretion rate. This indicates that the observed longitude variations are due to a changing location of the accretion spot and not a change in the relative orientation of the magnetic axis." + }, + "0112/astro-ph0112084_arXiv.txt": { + "abstract": "We report the detection of a extrasolar planet candidate orbiting the G1~V star HD 39091. The orbital period is 2083 d and the eccentricity is 0.62. With a minimum (M~sin~$i$) mass of 10.3 M$_{\\rm JUP}$ this object falls near the high mass end of the observed planet mass function, and may plausibly be a brown dwarf. Other characteristics of this system, including orbital eccentricity and metallicity, are typical of the well populated class of radial velocity planets in eccentric orbits around metal-rich stars. ", + "introduction": "Since the discovery of the first extra-solar planet (Mayor \\& Queloz 1995), planetary detections have been dominated by northern hemisphere search programmes -- most prolifically by the high precision velocity programmes at Lick (e.g., Butler et al. 1996) and Keck (e.g. Vogt et al. 2000) and lower precision programmes at OHP (Baranne et al. 1996), McDonald Observatory (Cochran et al. 1997), AFOE (Noyes et al. 1997), and programmes at La Silla (Kurster et al. 2000; Queloz et al. 2000). Of these programmes, which achieve precisions of $\\sim$ 10~m~s$^{-1}$, only the latter has access to the sky south of $\\sim -20\\deg$, and these achieve precisions of $\\sim$ 10~m~s$^{-1}$. In 1998, the Anglo-Australian Planet Search (AAPS) began in the southern hemisphere enabling all-sky coverage of the brightest stars at precisions reaching 3~m~s$^{-1}$. In this paper we present the second set of results from this programme. Two companion papers, Tinney et al. (2002a, 2002b), present results for a further two planets and an initial investigation of the Calcium H and K activity among some of the target stars of the AAPS. Precision Doppler surveys have found all of the known extrasolar planets around solar-type stars. Discoveries have included: the first system of multiple planets orbiting a Sun-like star (Butler et al. 1999); the first planet seen in transit (Henry et al. 2000, Charbonneau et al. 2000); the first two sub-Saturn-mass planets (Marcy, Butler \\& Vogt 2000); and the AAPS' discovery of the first planet in a circular orbit outside the 0.1~au tidal-circularisation radius (Butler et al. 2001). A number of major surprises have emerged from the sample of extrasolar planets. \\begin{itemize} \\item The sub-stellar companion mass function for F, G, and K dwarfs rises strongly below 10 M$_{\\rm JUP}$, and shows no signs of flattening toward the detection limit near 1 M$_{\\rm JUP}$. And surprisingly, brown dwarf companions to solar type stars are rare (Butler et al. 2000). This is in spite of a strong selection bias in the observations that makes brown dwarfs easier to detect than planets. It should be noted that the mass is only known within the uncertainty of the projection factor represented by the sin~$i$ of the orbit. However, the sin~$i$ statistic is between 1 and 0.5 for 87 per cent of randomly inclined orbits, so 87 per cent of reported $M$ sin $i$ values will be within a factor of two of the true masses. \\item About $\\sim$0.75\\% of nearby solar-type stars stars have been found to have planets orbiting in circularised orbits inward of 0.1~au (the 51\\,Peg-like ``hot Jupiters''). A smaller fraction of stars is now being found to have ``eps Ret -like'' planets orbiting in circularised orbits at 1 au or so (Tinney et al. 2002). But the dominant class of extra solar planets, (found around some 7 per cent of target stars) show highly eccentric orbits within 3.5 au. None of these classes were predicted {\\em a priori} by planetary formation theories. Such theories have received enormous impetus from these observations, leading to models for planet formation and evolution which now include the effects of dynamical friction, disk-planet and planet-planet interactions (e.g., Rasio \\& Ford 1996; Weidenschilling \\& Marzari 1996; Artymowicz 2000; Boss 2000). \\item The majority of extrasolar planets that have been found so far occur around stars which are mildy metal-rich ($\\sim$+0.2 dex with considerable scatter) compared to the Sun (Laughlin 2000; Santos, Israelian \\& Mayor 2002). \\end{itemize} This paper reports the discovery of a new planet candidate from the AAPS. Section 2 describes this precision programme. The stellar properties and Keplerian orbital fits for the new planet candidate are presented in Section 3. Section 4 provides a discussion of the new object. ", + "conclusions": "We present data showing evidence for an eccentric-orbit extra-solar planet around the metal-rich star HD~39091. The detection of this long period object gives added impetus for the continuation of these searches to longer periods. We now must endeavour to continue to improve the precision and stability of the AAPS to be sensitive to the 10+ year periods where analogues of the gas giants in our own Solar System may become detectable around other stars." + }, + "0112/astro-ph0112421_arXiv.txt": { + "abstract": "Mergers of massive gas-rich galaxies trigger violent starbursts that -- over timescales of $> 100$ Myr and regions $> 10$ kpc -- form massive and compact star clusters comparable in mass and radii to Galactic globular clusters. The star formation efficiency is higher by 1 -- 2 orders of magnitude in these bursts than in undisturbed spirals, irregulars or even BCDs. We ask the question if star formation in these extreme regimes is just a scaled-up version of the normal star formation mode of if the formation of globular clusters reveals fundamentally different conditions. ", + "introduction": "The basic question I'm going to address is if {\\bf S}tar {\\bf F}ormation ({\\bf SF}) is a universal process just scaled up in intensity in the violent starbursts going on in merging gas-rich massive galaxies from what it is, at a much lower level, in normal spiral, dwarf, or low surface brightness galaxies, or if these vastly different sites feature different modes of SF. \\medskip\\noindent Normal SF in undisturbed galaxies produces stars, OB-associations, open clusters with masses up to ${\\rm 10^3~M_{\\odot}}$, and, sometimes, even a couple of Super Star Clusters (Larsen \\& Richtler 1999). It does, however, not seem to produce {\\bf G}lobular {\\bf C}lusters (at least not in large numbers). Is GC formation a process different from that for open cluster/association/field star formation or is there {\\bf one} universal process forming the entire continuum from associations through GCs? The answer may come from analyses of the molecular cloud structure and mass spectra or of the {\\bf M}ass {\\bf F}unction ({\\bf MF}) of young star clusters, globular and open. \\medskip\\noindent Violent SF occurs in interacting and merging gas-rich galaxies with the most powerful starbursts going on in Ultraluminous Infrared Galaxies with SF rates of several ${\\rm 100~M_{\\odot}yr^{-1}}$. Hydrodynamic modelling has shown that high gas pressure is built up in the inner parts of massive merging galaxies, observations have shown that sometimes enormous amounts of molecular gas at extremely high desities ${\\rm > 1000~M_{\\odot}pc^{-3}}$, comparable to the stellar densities in the cores of elliptical galaxies, are assembled in regions of 1 -- few kpc extent. SF efficiencies are expected to be very high there. Indeed, the strong Balmer absorption lines and the still high luminosity of NGC 7252, a $\\lta 1$ Gyr old merger remnant, under conservative assumptions concerning gas content G and luminosity of the progenitor spirals imply that over a scale of order 10 kpc the global starburst in this galaxy must have had a SF efficiency ${\\rm \\eta := \\frac{\\Delta S}{G}}$ of order 30 \\% or more. This is at least a factor 10 higher than what is observed in spirals, irregulars, and BCDGs. It is high enough to allow for the formation of massive compact star clusters that may survive over a Hubble time, i.e. young GCs (Fritze -- v. A. \\& Gerhard 1994, Fritze -- v. A. \\& Burkert 1995). Is the formation of GCs a criterium that discriminates between violent and normal SF regimes? What fraction of the Super Star Clusters observed in interacting galaxies are young GCs? ", + "conclusions": "" + }, + "0112/astro-ph0112435_arXiv.txt": { + "abstract": "We present the initial results from an [O~III] $\\lambda$ 5007 survey for intra-group planetary nebulae in the M~81 group of galaxies. A total of 0.36 square degrees of the survey have been analyzed thus far, and a total of four intra-group candidates have been detected. These data allow us to probe the physics of galaxy interactions in small groups, and give us an upper limit for the density of intracluster starlight. We find that the M~81 group has less than 3\\% of its stars in an intra-group component; this is much less than the fraction seen in richer galaxy clusters. ", + "introduction": " ", + "conclusions": "" + }, + "0112/astro-ph0112059_arXiv.txt": { + "abstract": "While the modern stellar IMF shows a rapid decline with increasing mass, theoretical investigations suggest that very massive stars ($\\gtrsim100\\,\\Msun$) may have been abundant in the early universe. Other calculations also indicate that, lacking metals, these same stars reach their late evolutionary stages without appreciable mass loss. After central helium burning, they encounter the electron-positron pair instability, collapse, and burn oxygen and silicon explosively. If sufficient energy is released by the burning, these stars explode as brilliant supernovae with energies up to 100 times that of an ordinary core collapse supernova. They also eject up to 50\\,\\Msun of radioactive $^{56}$Ni. Stars less massive than 140\\,\\Msun or more massive than 260\\,\\Msun should collapse into black holes instead of exploding, thus bounding the pair-creation supernovae with regions of stellar mass that are nucleosynthetically sterile. Pair-instability supernovae might be detectable in the near infrared out to redshifts of 20 or more and their ashes should leave a distinctive nucleosynthetic pattern. ", + "introduction": "Owing to the lack of any metals, the cooling processes that govern star formation are greatly reduced for first generation of stars (Pop III). Magnetic fields and turbulence may also be less important at these early times \\cite{ABN00}. Consequently, theoretical studies \\cite{Lar98} indicate that the Jeans mass for primordial stars in their special environment may have been as great as $\\sim1000\\,\\Msun$. Numerical simulation of primordial star formation predict the occurrence of such stars at red shifts $\\sim20$ and an initial mass function (IMF) that either peaks at $\\sim100\\,\\Msun$ \\cite{ABN00,BCL99} or is bimodal \\cite{NU00}, i.e., also contains stars of a few \\Msun. Once formed, at solar metallicity, massive stars ordinarily experience significant mass loss \\cite{Fig98} and may end as relatively small objects, but for low metallicity mass loss is suppressed. In \\S~\\ref{loss} we discuss the peculiarities of mass loss and evolution of very massive primordial stars. Figure~\\ref{MM3} gives an overview of expected final fates of metal-free stars as a function of initial mass. In \\S~\\ref{z20} we examine the expected light curve of a pair-creation supernova from a 250\\,\\Msun star at a redshift of $z=20$ and in \\S~\\ref{nuc} we review nucleosynthetic yields from Pop III. Some conclusions are given in \\S~\\ref{end}. \\begin{figure} \\includegraphics[angle=270,width=\\columnwidth]{hegeraf1.ps} \\caption{Initial-final mass function of non-rotating Pop III stars. The \\textsl{x-axis} gives the initial mass. The \\textsl{y-axis} gives both the final mass of the collapsed remnant (\\textsl{thick black curve}) and the mass of the star when the event begins that produces that remnant (e.g., mass loss in AGB stars, supernova explosion for those stars that make a neutron star, etc.; \\textsl{thick gray curve}). We distinguish four regimes of initial mass: \\emph{low mass stars} below $\\sim10\\,\\Msun$ that end as white dwarfs; \\emph{massive stars} between $\\sim10\\,\\Msun$ and $\\sim100\\,\\Msun$ that form an iron core that eventually collapses; \\emph{very massive stars} between $\\sim100\\,\\Msun$ and $\\sim1000\\,\\Msun$ that encounter the pair instability; and \\emph{supermassive stars} (arbitrarily) above $\\sim1000\\,\\Msun$. Since no mass loss is expected for $Z=0$ stars before the final stage, the grey curve is approximately the same as the line of no mass loss (\\textsl{dotted}). Exceptions are $\\sim100-140\\,\\Msun$ where the pulsational pair instability ejects the outer layers of the star before it collapses, and above $\\sim500\\,\\Msun$ where pulsational instabilities in red supergiants may lead to significant mass loss \\cite{BHW01}. Since the magnitude of the latter is uncertain, lines are drawn \\textsl{dashed}. For a more detailed description, please refer to \\cite{HW01}\\label{MM3}} \\end{figure} ", + "conclusions": "\\label{end} Due to the absence of metals the first generation of star will likely not experience significant mass loss by radiation-driven stellar winds or opacity-driven pulsations. Their unique structure also prevents significant mass loss by the epsilon mechanism. Therefore very massive Pop III single stars reach carbon burning with enough mass to encounter the pair instability. Since current theoretical studies indicate that such star may constitute a significant, if not dominant, fraction of Pop III, we predict that their nucleosynthetic yields may have a unique imprint on the chemical evolution of the early universe. Their production of odd-$Z$ elements is by $\\sim2$ orders of magnitude lower than that of even-$Z$. Elements heavier than zinc are not produced. Since stars of lower mass ($\\lesssim 140\\,\\Msun$) or higher mass ($\\gtrsim 260\\,\\Msun$) collapse into black holes without significant heavy element creation. Pair-creation SNe are thus a ``clean'' source of nucleosynthesis in the sense that neighboring mass ranges do not ``pollute'' the sample. In case of a bimodal Pop III IMF, massive stars in the range $\\sim8\\ldots40\\,\\Msun$ will also contribute to the resulting abundance pattern, though on a slightly slower time-scale (factor $\\sim2$). Also in them, the lack of initial CNO ``seeds'' will lead to an elemental odd-even pattern, though much less expressed \\cite{HW01}, and they possibly contribute $r$-process isotopes. It is even conceivable that Pop III AGB stars contribute some $s$-process \\cite{ven01}. The interaction of the ejecta with the surrounding matter, a possible enrichment of the intergalactic medium and mixing of contributions from different-mass sources before the formation of the first Pop II stars will have to be studied in more detail in the future. We predict that Pop III pair-creation SNe might be detectable by future near infrared space experiments --- all the way out to the edge of universe --- to redshifts of 20 or more. Combined with the challenge to find old Pop II stars that show the predicted abundance pattern from the ashes of these explosions, this should allow deeper insight into the happenings at the times when the first sparks of stellar light terminated the ``dark ages''. {\\footnotesize \\textbf" + }, + "0112/astro-ph0112090_arXiv.txt": { + "abstract": "The transfer of polarized radiation in stochastic synchrotron sources is explored by means of analytic treatment and Monte Carlo simulations. We argue that the main mechanism responsible for the circular polarization properties of compact synchrotron sources is likely to be Faraday conversion and that, contrary to common expectation, a significant rate of Faraday rotation does not necessarily imply strong depolarization. The long-term persistence of the sign of circular polarization, observed in many sources, is most likely due to a small net magnetic flux generated in the central engine, carried along the jet axis and superimposed on a highly turbulent magnetic field. We show that the mean levels of circular and linear polarizations depend on the number of field reversals along the line of sight and that the gradient in Faraday rotation across turbulent regions can lead to ``correlation depolarization''. Our model is potentially applicable to a wide range of synchrotron sources. In particular, we demonstrate how our model can naturally explain the excess of circular over linear polarization in the Galactic Center and the nearby spiral galaxy M81 and discuss its application to the quasar 3C 279, the intraday variable blazar PKS 1519-273 and the X-ray binary SS 433. ", + "introduction": "Polarization has proven to be an important tool in AGN research. In principle, linear and particularly circular polarization observations of synchrotron radiation may permit measurements of various properties of jets such as: magnetic field strength and topology, the net magnetic flux carried by jets (and hence generated in the central engine), the energy spectrum of radiating particles, and the jet composition (i.e., whether jets are mainly composed of $e^{+}-e^{-}$ pairs or electron-proton plasma). The renewed interest in polarization of compact radio sources stems from two recent developments. First, \\citet{bow99} detected circular polarization using the Very Large Array (VLA) in the best supermassive black hole candidate, the Galactic center (Sgr A$^{*}$). This discovery was quickly confirmed by \\citet{sau99} using the Australia Telescope Compact Array (ATCA). Circular polarization was also detected in the celebrated X-ray binary system SS 433 \\citep{fen00}. Moreover, the Very Long Baseline Array (VLBA) has now detected circular polarization in as many as 20 AGN \\citep{war98,hom99}. Second, it is now possible to measure circular polarization with unprecedented accuracy of 0.01\\% using the ATCA \\citep{ray00}. This dramatic improvement in the observational status of polarization measurements has also brought new questions. For example, there is now growing observational evidence that the sign of circular polarization is persistent over decades \\citep{kom84,hom99}, which indicates that it is a fundamental property of jets. Another problem, which has not been satisfactorily explained as yet, is how to reconcile the high level of circular polarization with the lower value of linear polarization in Sgr A$^{*}$ \\citep{bow99} and M81$^{*}$ \\citep{bru01}. Indeed, there is not even a general consensus on the mechanism responsible for the circular polarization properties of jets \\citep{war98}. \\\\ \\indent In this paper we attempt to solve some of the theoretical puzzles. The paper is organized as follows. In the next section we summarize the most important observational facts. In Section 3 we briefly discuss mechanisms for producing circular polarization and argue that the most likely process is ``Faraday conversion''. Section 4 presents our model for polarization and in the subsequent sections we compare Monte Carlo simulations with analytic results and discuss general implications for observations as well as specific observational cases. We summarize our conclusions in Section 7. ", + "conclusions": "We have considered the transfer of polarized synchrotron radiation in stochastic sources by means of an analytic approach and a set of numerical simulations, and have argued that Faraday conversion is the primary mechanism responsible for the circular polarization properties of compact radio sources. A crucial ingredient of our model is a small bias in the highly turbulent magnetic field which accounts for the persistence of the sign of circular polarization. This bias is direct evidence for the net magnetic flux carried by magnetically accelerated jets (e.g., \\citet{blap82,lic92}).\\\\ \\indent Extremely large rates of Faraday rotation, i.e., Faraday rotation per unit synchrotron absorption depth, do not necessarily lead to depolarization provided that the mean rate of Faraday rotation across the source is relatively small, or in other words, that the turbulent magnetic field possesses a very small directional bias. Indeed, a large Faraday rotativity is required in order to explain the high ratio of circular to linear polarization observed in some sources. Constraints on jet composition or accretion rate, based on the requirement that the source does not become Faraday depolarized, may be circumvented under these conditions.\\\\ \\indent Gradients in Faraday rotation across turbulent cells can lead to correlations between rotativity and Stokes $Q$ and $U$ parameters, which can result in ``correlation depolarization''. Observed polarization levels require that the field have many reversals along the line of sight to avoid this effect. Statistical fluctuations of circular and linear polarizations are then likely to be dominated by changes in the mean parameters describing the plasma rather than by the stochastic behavior of the turbulent medium. Variations in the mean parameters are unlikely to change the helicity of circular polarization unless a source undergoes a sharp transition from very low to very high synchrotron depth.\\\\ \\indent We have shown that our model is potentially applicable to a wide range of compact synchrotron sources. In particular, it naturally predicts an excess of circular over linear polarization when a source is strongly depolarized by the mean Faraday rotation and when a small amount of linear polarizarion is efficiently converted into circular polarization. This can explain the polarization properties of the Galactic Center and M81$^{*}$." + }, + "0112/astro-ph0112215_arXiv.txt": { + "abstract": "An {\\it XMM-Newton} observation of the luminous Seyfert 1 galaxy 1H~0419-577 is presented. We find that the spectrum is well fitted by a power law of canonical slope ($\\Gamma$~$\\sim 1.9$) and 3 blackbody components (to model the strong soft excess). The {\\it XMM} data are compared and contrasted with observations by {\\it ROSAT} in 1992 and by {\\it ASCA} and {\\it BeppoSAX} in 1996. We find that the overall X-ray spectrum has changed substantially over the period, and suggest that the changes are driven by the soft X-ray component. When bright, as in our {\\it XMM-Newton} observation, it appears that the enhanced soft flux cools the Comptonising corona, causing the 2--10 keV power law to assume a `typical' slope, in contrast to the unusually hard (`photon-starved') spectra observed by {\\it ASCA} and {\\it BeppoSAX} four years earlier. ", + "introduction": "\\label{sec:intro} 1H~0419-577 (also known as LB~1727, 1ES~0425-573 and IRAS~F04250-5718) is a radio-quiet Seyfert galaxy, with a $60 \\mu m$ flux of 0.18~Jy and an apparent magnitude of 14.1. It is a moderate redshift object ({\\it z} $=0.104$) and relatively bright X-ray source which has been observed over recent years by {\\it ASCA}, {\\it ROSAT} and {\\it BeppoSAX}. 1H~0419-577 was also one of the brightest Seyfert galaxies detected in the extreme ultra-violet by the {\\it ROSAT} Wide Field Camera (Pye \\etal 1995) and {\\it EUVE} (Marshall, Fruscione \\& Carone 1995). Optical spectra taken over the same period in 1996 as the {\\it ASCA} and {\\it BeppoSAX} X-ray observations (Guainazzi \\etal 1998) show 1H~0419-577 to be a typical broad-line Seyfert 1, in accordance with the classification by Brissenden (1989). Over the 2--10~keV band, Seyfert galaxies can usually be modelled by a power law, with photon-index $\\Gamma$ $\\sim 1.8-2$. Below about 1 keV a `soft excess' is often reported, although the limited bandwidth and resolution of previous missions have made it difficult to distinguish a soft {\\it emission} component from the effects of absorption by ionised matter. In the case of 1H~0419-577, Turner \\etal (1999) did conclude that there is a soft emission component on the basis of simultaneous {\\it ROSAT} HRI and {\\it ASCA} observations. However, those authors found the 2--10~keV power law to be unusually flat, with $\\Gamma$ $\\sim 1.5-1.6$, and to extend down to 0.7~keV. The {\\it BeppoSAX} observation of 1H~0419$-$577 in 1996 September (Guainazzi \\etal 1998) also found an unusually flat power law (over 3--10 keV) of $\\Gamma$ $\\sim 1.55$, but provided no independent soft X-ray data, due to technical problems with the Low Energy Concentrator Spectrometer (LECS). The link between the hard X-ray power law and a soft X-ray emission component is critical in the context of the accretion disc/corona model for AGN, where the hard X-ray emission is explained by Comptonisation of optical/EUV photons from the disc by energetic electrons in an overlying corona (e.g., Haardt \\& Maraschi 1991). In this model the energy balance between the soft photon flux and the corona then determines the hardness of the spectrum in the 2--10~keV band. The soft excess may be the tail of the Big Blue Bump, representing the thermal emission from an accretion disc surrounding the central black hole (Shields 1979; Czerny \\& Elvis 1987; Ross \\& Fabian 1993). In our {\\it XMM-Newton} observation, reported here, we find a strong and broad soft emission component (consistent with an earlier {\\it ROSAT} PSPC observation in 1992; Guainazzi \\etal 1998) and a 2--10~keV power law continuum slope typical of Seyfert 1 galaxies. We discuss our result in terms of the stronger soft photon flux cooling the coronal electrons, with a resulting steepening of the power law, and suggest it represents direct observational support for the disc/corona model for the hard X-ray emission from radio-quiet AGN. In that general class of models, back-irradiation of the accretion disc by the hard X-ray flux can result in additional features being imprinted on the emerging X-ray spectrum. The most obvious of these `reflection' features (Pounds \\etal 1990; Nandra \\& Pounds 1994) is often an emission line at 6.4~keV arising from fluorescence in near-neutral Fe. This line has emerged as a powerful diagnostic of the inner regions in AGN since {\\it ASCA} observations found it to be broadened and red-shifted (Tanaka \\etal 1995; Nandra \\etal 1997a, b). Early observations from {\\it XMM-Newton} have shown a rather different situation to apply in several high-luminosity Seyferts (similar to 1H~0419-577), with a weaker and higher energy (ionised) broad Fe-K line, resolved from a narrow line at 6.4~keV (eg. Reeves \\etal 2001; Pounds \\etal 2001). The latter component, interpreted as scattering from neutral matter distant from the hard X-ray source (e.g., in the molecular torus), is emerging as a common feature in AGN observations by {\\it XMM-Newton} and {\\it Chandra}. In previous observations of 1H~0419-577 Turner \\etal (1999) found evidence for an emission line in the {\\it ASCA} data at 6.39~keV, with equivalent width (EW) of 700~$\\pm$~400~eV (data from August 1996). A line is not detected in the {\\it SAX} data (Guainazzi \\etal 1998), but with a rather high upper limit of $\\sim$~250~eV for the equivalent width. In Section~\\ref{sec:xmmobs} the data from {\\it XMM-Newton} are summarised, followed by the data analysis in Section~\\ref{sec:specanal}. Comparisons with the {\\it ROSAT}, {\\it ASCA} and {\\it BeppoSAX} observations are reviewed in Section~\\ref{sec:otherobs}. Note that all fit parameters are given for the rest frame of the AGN, with values of $H_0$~=~50 km~s$^{-1}$~Mpc$^{-1}$ and $q_0$~=~0 assumed throughout. Errors are quoted at the 90\\% confidence level ($\\Delta \\chi^{2}$~=2.7 for 1 interesting parameter). ", + "conclusions": "A new measurement of the X-ray spectrum of the luminous Seyfert 1 galaxy 1H 0419-577 with {\\it XMM-Newton} has shown it to have a strong soft excess and `normal' power law slope, very similar to the comparably luminous Seyferts Mrk 205 and 509. No iron line was found, but the upper limits are consistent with the Fe-K emission found in {\\it XMM-Newton} observations of objects of similar luminosity. A comparison with archival data shows the 2--10~keV continuum to have been much flatter in 1996, together with a far weaker soft excess. We interpret our overall {\\it XMM-Newton} observation in terms of two-temperature Comptonisation of thermal photons from the inner accretion disc, and explain the hard continuum seen in 1996 as being due to `photon starving' of the hotter electron component. We note the spectral changes we have observed in the Seyfert galaxy 1H 0419-577 are similar to, but smaller than, those seen in the changes in `state' for several galactic black hole candidate sources (e.g., Vilhu \\etal 2001)." + }, + "0112/astro-ph0112509_arXiv.txt": { + "abstract": "{We study magnetically powered relativistic outflows in which a part of the magnetic energy is dissipated internally by reconnection. For GRB parameters, and assuming that the reconnection speed scales with the Alfv\\'en speed, significant dissipation can take place both inside and outside the photosphere of the flow. The process leads to a steady increase of the flow Lorentz factor with radius. With an analytic model we show how the efficiency of this process depends on GRB parameters. Estimates are given for the thermal and non-thermal radiation expected to be emitted from the photosphere and the optically thin part of the flow respectively. A critical parameter of the model is the ratio of Poynting flux to kinetic energy flux at some initial radius of the flow. For a large value ($\\ga 100$) the non-thermal radiation dominates over the thermal component. If the ratio is small ($\\la 40$) only prompt thermal emission is expected which can be identified with X-ray flashes. ", + "introduction": "\\label{sec:introduction} To overcome the compactness problem of $\\gamma$-ray bursts (GRBs) \\citep[e.g.][]{piran:99} the central engines must produce radiating material moving ultra-relativistically fast towards the observer. GRB models must therefore describe an energy source which not only releases energy of around $10^{52}\\,\\mathrm{erg/sterad}$ but must also explain the `clean' form of the energy. To produce the high Lorentz factors of the order of $10^2$--$10^3$ \\citep{fenimore:93,woods:95,lithwick:01} which are needed only a small fraction of the total energy can exist in form of rest mass energy of the matter involved. The popular models involving compact objects or the collapse of a massive star to a black hole must include mechanisms how the energy is transported into a space region with few baryons. Otherwise large amounts of mass are expelled which cannot be accelerated to high Lorentz factors. The initially released energy could leave the central polluted region by neutrinos which annihilate to a pair plasma further away \\citep{berezinskii:87,goodman:87,ruffert:97}. But due to the small cross section of neutrinos the efficiency is low and most of the energy escapes as neutrinos. A Poynting flux dominating outflow will naturally occur if the compact object rotates and possesses a magnetic field. The luminosity will be fed by the rotational energy reservoir of the central object. Models involving an magnetised torus around a black hole \\citep{meszaros:97} or a highly magnetised millisecond pulsar \\citep{usov:92,kluzniak:98,spruit:99} would produce such a rotationally driven Poynting flux. Extraction of energy from the central object by this magnetic process is potentially very efficient and fast. In order to obtain not only a large energy extraction but also the observed large bulk Lorentz factors, the Poynting flux must be converted to kinetic energy. The simplest available magnetic acceleration models, in which the flow is approximated as radial, are problematic in this respect. In the classic non-relativistic case \\citep{weber:67,belcher:76} a dominating initial Poynting flux can transfer $1/3$ of its energy to the matter. If the flow is initially relativistic however almost no acceleration is possible \\citep{michel:69}. The physical reason lies in the singular field and flow geometry of a purely radial flow. In this case the magnetic pressure gradient balances the magnetic tension force and no acceleration occurs. An imbalance between the pressure gradient and the tension force occurs in non-radial outflows, if the flow lines diverge faster with radius than in the radial case \\citep{begelman:94,takahashi:98}. Detailed 1-dimensional calculations have been made which show how such a flow divergence can come about \\citep{beskin:97,daigne:01}. In this paper we show that there is a second process which naturally leads to efficient conversion of Poynting flux to bulk kinetic energy. If the magnetic field in the outflow contains changes of direction on sufficiently small scales, (a part of) the magnetic energy is `free energy' which can be released locally in the flow by `fast reconnection' processes. Such a decay of magnetic energy, if it can occur rapidly enough, has two desirable effects. First it provides a source of energy outside the photosphere which is converted directly into radiation, without the relatively inefficient intermediate step of internal shocks \\citep[ hereafter \\citetalias{spruit:01}]{spruit:01}. Secondly, it leads to an outward decrease of magnetic pressure, which causes a strong acceleration of the flow and conversion of Poynting flux to kinetic energy. In the present work, we concentrate on the acceleration effect, and show how it depends on the parameters (energy flux, baryon loading) of a GRB. This aspect of the model can be illustrated with analytic calculations. In a future paper we show, with more detailed numerical results, how the dissipated magnetic energy can also power the observed prompt radiation. Changes of direction of field lines must occur in the flow in order for energy release by reconnection to be possible. These can occur naturally in a number of ways. If the magnetic field of a rotating central object is \\emph{non-axisymmetric} the azimuthal part of the magnetic field in the flow changes direction on a length scale $\\lambda\\approx \\pi v/\\Omega$, where $v$ is the flow velocity and $\\Omega$ the angular frequency. For an inclined dipole this yields the `striped' field in pulsar wind model of \\citet{coroniti:90} where magnetic energy is released by the annihilation of the antiparallel field components. Field decay by reconnection was applied to pulsar winds \\citep{coroniti:90,lyubarsky:01} and also to GRBs (\\citealp{thompson:94}; \\citetalias{spruit:01}). In this paper we investigate the dynamics of a magnetically powered outflow in which some of the energy dissipates by reconnection. With the assumption that the flow is highly dominated by magnetic energy and that the thermal energy is negligible we derive the velocity profile of the flow. The results provide estimates of the Lorentz factor of the flow, the photospheric radius, and the amount of energy that can be converted into non-thermal radiation. We investigate under which conditions prompt emission is expected and whether a considerable amount of thermal radiation can be produced. These predictions can then be tested against observations of the thermal component in GRB spectra \\citep{preece:00}. ", + "conclusions": "\\label{sec:disc} We have investigated the effect of dissipation of magnetic energy in a GRB outflow on the acceleration of the flow. Such dissipation is expected if the flow contains small scale changes of direction of the field for example when the flow is produced by the the rotation of a non-axisymmetric magnetic field. The dissipation is governed by the speed of fast reconnection, parameterised in our calculations as a fraction $\\epsilon\\approx 0.1$ of the local Alfv\\'en speed in the flow. Two possibilities for the field geometry in the outflow have been considered: a geometry where the changes in the small scale field direction occur along the bulk flow direction, and a geometry where the field variation is transversal to the flow direction. The first mentioned, \\emph{longitudinal case} is expected in the equatorial plane of an inclined rotator as in the `striped' pulsar wind model of \\citet{coroniti:90}. The second, \\emph{transversal case} can be associated with a polar outflow where the field line structure resembles a spiral. In both cases there are MHD instabilities (tearing and kink instabilities) which lead to reconnection processes. They differ only by the functional form of the reconnection time scale. We find that in any case the process leads to a strong increase of the bulk Lorentz factor of the flow. This acceleration is due to the outward decrease of the magnetic pressure resulting from the field decay. At the same time, the dissipated energy can be released to large extend in the optically thin part of the flow beyond its photosphere, and can power most if not all of the prompt emission. This provides an alternative to the internal shock model. The calculation is done for a stationary wind. Why this approximation is valid for highly variable objects like GRBs is not obvious. The duration of GRBs $t$ is of the order of a few seconds. One can approximate the wind as stationary within a source distance $ct\\approx10^{11}\\,\\mathrm{cm}$. Thus the flow up to the photospheric radius is well described by a stationary description. Further out the time dependence of a real flow will become more important but that topic is beyond the scope of this work. The outflow with transversal field variation contains some additional complications which does not occur in the longitudinal case. The dissipation time scale is proportional to the source distance. This results in a rapid energy dissipation near to the source and the velocity profile depends critically on the radius where the dissipation sets in. But this initial radius is hard to estimate from first principles. We have used the spiral-like field geometry of a polar flow as pictured in \\citetalias{spruit:01} to justify the existence of transversal field variations. This field geometry occurs for a polar outflow of an axisymmetric rotator. The following arguments give reasons why this field geometry is rather special and may not be important in a general. The kink instability leads to a break-down of the ordered spiral field configuration. After some Alfv\\'en crossing times the field geometry will have changed so that the `longitudinal' dissipation time will become important while the `transversal' time scale grows large and can be neglected. On the other hand the rotator may not be perfectly aligned and non-axisymmetric field components are also present in the polar outflow. So, we probably have always longitudinal field variations in the flow so that the findings found in our treatment of the `longitudinal case' might be much more applicable and general. We assume that the thermal energy flux is negligible compared to the kinetic and Poynting energy flux. The temperature is set to zero which simplifies the treatment and allows an analytical integration of the dynamic equations. Setting the thermal pressure gradient artificially to zero might appear to underestimate the acceleration. On the other hand the energy equation takes care that all released magnetic energy shows up in kinetic form. In fact, we overestimate the acceleration by doing so because the energy part converted into heat reduces the the gain of kinetic energy in the flow. Another physical argument explains why the flow stays cold: The acceleration expressed by the scaling of the Lorentz factor gives $\\Gamma\\sim r^{1/3}$ for our model. The release of magnetic energy must therefore also scale with $r^{1/3}$. In contrast to that, purely thermal acceleration by adiabatic cooling leads to more rapid flow acceleration where the Lorentz factor scales like $\\Gamma\\propto r$ \\citep{paczynski:86}. Thus, heating proceeds slower than adiabatic cooling so that the thermal pressure gradient is not important compared to the magnetic pressure gradient which drives the flow. The reason why \\citet{lyubarsky:01} find a faster acceleration of $\\Gamma\\sim r^{1/2}$ in a similar model lies in the different reconnection prescription and is not due to their inclusion of thermal pressure. In the optically thin regime part of the dissipated energy radiates away. There, the model over-estimates the gain of kinetic energy. We cannot give arguments how much dissipated energy escapes as prompt radiation so that the total amount of released energy gives only an upper limit on the Lorentz factor. The photospheric radius determines the lower limit on radius for the region in which non-thermal radiation is expected to originate. For typical GRB parameters describing the total luminosity, the baryon loading, the fraction of dissipatable energy and the reconnection rate one finds that a considerable amount of dissipation takes place in the optically thin region. Part of the dissipated energy is converted into non-thermal radiation. The remainder still leads to an acceleration of the flow. This acceleration is caused by the magnetic pressure gradient induced by the field dissipation. Since the acceleration continues outside the photosphere up to the radius where all the free magnetic energy is used up this non-thermal radiation is emitted from matter with different Lorentz factors. The observable spectrum in thus smeared out compared to a spectrum from a uniformly moving medium. For a more sound analysis of this topic one needs a model for the radiation process. The Poynting flux conversion happens at radii $r\\la 10^{15}\\,\\mathrm{cm}$ which is inside the distance $\\approx10^{16}\\,\\mathrm{cm}$ where the GRB outflow is expected to run into the external medium. Thus, the Poynting flux can be converted efficiently. But by applying the model to the Crab pulsar wind we come to the same conclusions as \\citet{lyubarsky:01}: The conversion is inefficient since the observed pulsar wind bubble is to small to contain the whole region where reconnection takes place. For the Crab pulsar the assumed initial Poynting flux ratio is larger than for GRBs leading to a much longer reconnection phase. The presented model does not settle this Crab wind problem. The most important parameter which controls the amount of energy dissipated beyond the photosphere is the initial Poynting flux to kinetic energy flux ratio. If its value is around 100 or greater much non-thermal, prompt emission is produced. If its value is of the order of 10, however, all the Poynting flux energy is converted into kinetic energy and thermal radiation. Only prompt thermal emission and afterglow emission is expected in this case. The initial Poynting flux ratio is a measure for the baryon loading in a sense that a high baryon loading corresponds to a low initial Poynting flux ratio. Observations indicate that X-ray flashes and X-ray rich GRBs are very similar phenomena which probably differ only by the amount of baryon loading \\citep{heise:01}. In the context of our model, X-ray flashes can be associated with low initial Poynting flux ratios. In this case, the X-ray emission is thermal radiation from the photosphere. Increasing the initial Poynting flux ratio leads to the emission of non-thermal $\\gamma$-rays in the optically thin region, thus producing X-ray rich and regular GRBs. If afterglows of X-ray flashes could be observed they would yield information about the connection to regular GRBs. Afterglows depend less strongly on the initial Poynting flux ratio but rather on the total luminosity of the outflow. Thus, X-ray flash afterglows should be similar to afterglows of regular GRBs according to our model. In a future work we will investigate the thermal emission more quantitatively. The model predicts black-body radiation originating from the photosphere of the flow. We can calculate the radius of the photosphere and the Lorentz factor of the flow there. Together with the temperature one is able to calculate the luminosity if the thermal radiation. Since our approximation treats the flow as cold we cannot give quantitative results in this respect. Though, one finds that the Lorentz factor at the photosphere depends only weakly on the model parameters. Therefore, the observable temperature $kT_\\mathrm{obs} = \\Gamma_\\mathrm{ph} kT/(1+z)$ of the thermal component of a GRB depends primarily on the redshift $z$ and the temperature in the comoving frame $T$. This result simplifies the task to disentangle the effects of different model parameters on the temperature. A detailed, quantitative analysis of the thermal radiation will be done in a following study." + }, + "0112/astro-ph0112353_arXiv.txt": { + "abstract": "We discuss three classes of x-ray transients to highlight three new types of transients found with the Wide Field Cameras onboard BeppoSAX. First there are the transients related to Low Mass X-ray Binaries in outburst, typically lasting weeks to months and reaching luminosities of the Eddington limit for a few solar masses. Recently another subclass of outbursts in such binaries has been discovered, which are an order of magnitude fainter and last shorter than typical hours to days. We discuss whether they constitute a separate subset of x-ray binaries. A second class of x-ray transients are the x-ray bursts. Thermonuclear explosions on a neutron star (type~I x-ray bursts) usually last of order minutes or less. We discovered a second type (called super x-ray bursts) with a duration of several hours. They relate to thermonuclear detonations much deeper in the neutron star atmosphere, possibly burning on the nuclear ashes of normal x-ray bursts. The third class are the enigmatic Fast X-ray Transients occurring at all galactic latitudes. We found that the bright ones are of two types only: either nearby coronal sources (lasting hours) or the socalled x-ray flashes (lasting minutes). The new class, the X-ray flashes, may be a new type of cosmic explosion, intermediate between supernovae and gamma ray bursts, or they may be highly redshifted gamma ray bursts. It thus appears that the three classes of x-ray transients each come in two flavors: long and short. ", + "introduction": "\\begin{figure}[h] \\centering \\hbox{ \\psfig{file=heise_fig1.eps,width=0.45\\textwidth,clip=t} \\parbox[b]{.5\\textwidth}{\\caption[] {A ty\\-pi\\-cal WFC ima\\-ge of $40^o\\times 40^o$ a\\-round the ga\\-lac\\-tic cen\\-ter. The ang\\-ul\\-ar re\\-so\\-lu\\-tion is 5\\arcmin. This field con\\-tains about $50\\%$ of all LMXBs in the gal\\-axy. Po\\-si\\-tion de\\-ter\\-mi\\-na\\-tions have an ac\\-cu\\-ra\\-cy of ty\\-pi\\-cal\\-ly 1\\arcmin -3\\arcmin. About 30 point sour\\-ces are seen ab\\-ove the de\\-tec\\-tion thres\\-hold of ty\\-pi\\-cal\\-ly 10 mCrab in $10^5$ s. \\vspace*{3mm} }% }% }% \\label{wfcGC} \\end{figure} The x-ray sky is transient in nature. X-ray sources appear and disappear. This fact has been known since the beginning of x-ray astronomy, soon after the discovery of the first few x-ray sources. Before space born instruments, x-ray detectors were placed on sounding rockets and scanned the entire sky in a spinning payload during a few minutes above the atmosphere. On the time scale of successive launches (months to years) bright x-ray sources may vanish below detection limits and new sources may be among the brightest x-ray sources. Later, variablity on all time scales (millisec to years) became a well known characteristic of the x-ray sky. For x-ray binaries, many of these time scales are now explained in terms of accretion onto compact objects. In this review I limit the discussion on these transients to two, recently discovered, phenomena: faint transients and x-ray superbursts. An entirely different class of transient phenomena are the Fast X-ray Transients (FXTs) or High Latitude Transients. They were discovered (Forman et al. 1978, Pye \\& McHardy 1983, Ambruster et al. 1986) with the first few satellite born x-ray instruments, such as those on board UHURU, Vela, Ariel V, HEAO-1, etc. These rotating satellites scanned the entire sky. The scan frequency often occurred on the same time scale as the satellite orbit, typically 1.5 hours. Such events are detected in one sky scan and disappeared soon afterward, typically limiting the duration to be longer then a minute and shorter than a few hours. In some cases the FXT was detected on the next sky scan as well, setting a duration of order one hour. Because of the limited amount of data, the nature of these sources, although seen by almost every x-ray satellite, remains an enigma. The attitude mode and orbit of the satellite continues to determine the character of the time scale of x-ray sources. 3-axis stabilized x-ray satellites made the discovery possible of short and sparsely distributed outbursts in known sources. The source monitoring is only interrupted by earth occultations. The first of such a satellite was the Astronomical Netherlands Satellite ANS. With the x-ray instruments on ANS the first x-ray burst was discovered in 2U1820-30 (Grindlay \\& Heise 1975, Grindlay et al. 1976). They later later appeared to be thermonuclear explosions in the accreting layers of neutron stars in x-ray binaries. We discuss the recently discovered new type of thermonuclear explocions in section 3. ", + "conclusions": "" + }, + "0112/astro-ph0112486_arXiv.txt": { + "abstract": "{ Primordial black holes ({\\sc pbh}s) have motivated many studies since it was shown that they should evaporate and produce all kinds of particles (Hawking \\cite{Hawking1}). Recent experimental measurements of cosmic rays with great accuracy, theoretical investigations on the possible formation mechanisms and detailed evaporation processes have revived the interest in such astrophysical objects. This article aims at using the latest developments on antiproton propagation models (Maurin et al. \\cite{David} and Donato et al. \\cite{Fiorenza}) together with new data from BESS, CAPRICE and AMS experiments to constrain the local amount of {\\sc pbh} dark matter. Depending on the diffusion halo parameters and on the details of emission mechanisms, we derive an average upper limit of the order of $\\rho_{\\odot}^{PBH}\\approx 1.7\\cdot10^{-33}$g cm$^{-3}$. ", + "introduction": " ", + "conclusions": "\\subsection{Comparison with other existing limits} Many upper limits on the {\\sc pbh} explosion rate have been derived thanks to 100 MeV, 1 TeV and 100 TeV $\\gamma$-rays or low-energy antiprotons. In the ultra-high energy range, a reliable search for short gamma-ray bursts radiation from an arbitrary direction have been performed using the CYGNUS air-shower array (Alexandreas {\\it et al.} \\cite{Alex}). No strong one second burst was observed and the resulting upper limit is $dN_{\\odot exp}^{PBH}/dt\\, d^3V \\leq 0.9\\cdot10^6~{\\rm year}^{-1}{\\rm pc}^{-3}$. Very similar results were derived by the Tibet (Amenomori {\\it et al.} \\cite{Ame}) and the AIROBIC collaborations (Funk {\\it et al.} \\cite{Funk}). TeV gamma-rays have also been used to search for short time-scale coincidence events. The bursts detected are compatible with the expected background and the resulting upper limit obtained with 5 years of data (Connaughton \\cite{Connaughton}) is $dN_{\\odot exp}^{PBH}/dt\\, d^3V\\leq 3\\cdot10^6~{\\rm year}^{-1}{\\rm pc}^{-3}$. The limit coming from antiprotons has been advocated to be far better: the previous study from Maki {\\it et al.} (\\cite{Orito}) gives $dN_{\\odot exp}^{PBH}/dt\\, d^3V \\leq 2\\cdot10^{-2}~{\\rm year}^{-1}{\\rm pc}^{-3}$. However, we emphasize that this limit does not take into account the wide range of possible astrophysical uncertainties (in particular $L$, which can affect the limits by one order of magnitude). Moreover, we believe that the explosion rate is not the pertinent variable to use when comparing results from different approaches: thresholds differences between experiments make the meaning of \"explosion\" very different. With a mass spectrum $\\propto M^2$ for small masses, the number of exploding {\\sc pbh}s depends strongly on the value of the threshold. It makes the comparison of our results with the ones from Maki {\\it et al.} very ambiguous. This is why we prefer to give our upper limit, either as a local mass density, assuming a {\\it standard} mass spectrum, $\\rho^{PBH}_{\\odot} < 5.3\\cdot 10^{-33}~{\\rm g} \\,{\\rm cm}^{-3}$, either as a number density, independant of the mass spectrum shape outside the relevant interval, $n_{\\odot}^{PBH} < 1.3 \\cdot 10^{-50}~{\\rm cm}^{-3}$ (whatever L). Gamma-rays in the 100 MeV region provide a sensitive probe to the presence of PBHs along the line of sight up to redshifts as large as $\\sim 700$. Gamma-rays in this energy range have little interactions with the intergalactic medium and can travel on cosmological distances. The integration of the signal involves therefore a much larger scale than in the case of Milky Way antiprotons. It should also be pointed out that in this case most of the PBH population is involved as the dominant emission peaks above 100 MeV even at the present epoch. By matching the PBH cosmological emission to the extra-galactic gamma-ray diffuse background (MacGibbon \\& Carr \\cite{MacGibbon2}), the limit $\\Omega_{PBH} \\leq 1.8 \\times 10^{-8} \\, h^{-2}$ was obtained. Being mostly based on the assumption that a standard PBH mass spectrum holds above $M_*$, this result is robust. Our limit ($\\rho^{PBH}_{\\odot} < 5.3 \\times 10^{-33}~{\\rm g} \\,{\\rm cm}^{-3}$) also requires the same assumption. It no longer depends on the details of cosmic-ray propagation as it corresponds to the minimal possible value of 1 kpc for $L$. This result is therefore quite conservative. In order to discuss it in the light of the gamma-ray constraints, it should be noticed that any PBH population is a particular form of cold dark matter. When the latter collapses to form galactic halos and the intra-cluster medium, PBHs merely follow the collapse. Their abundance in the solar neighbourhood should trace their cosmological contribution to the overall value of $\\Omega_M \\sim 0.3 - 0.4$. Because a canonical isothermal halo has a solar density of $\\sim 0.3$ GeV cm$^{-3}$ -- $\\sim 5.3 \\times 10^{-25}$ g cm$^{-3}$ -- we infer an upper limit of $\\sim 10^{-8}$ on the contribution of PBHs to the galactic -- and as mentionned above to the cosmological -- dark matter. Our antiproton bound translates into $\\Omega_{PBH} \\leq 10^{-8} \\, \\Omega_M \\sim 4 \\times 10^{-9}$. Such a constraint is comparable to the limit derived from gamma-ray considerations. Antiprotons are produced by PBHs that are exploding at the present epoch and the limit which they provide is complementary to the gamma-ray constraint. \\subsection{Possible improvements on the antiproton limit} New data on stable nuclei along with a better understanding of diffusion models (see for example Donato et al., in preparation) could allow refinement in propagation to constrain $L$. It is also important to notice that the colour emission treatment could probably be improved. Studying more accurately the colour field confinement and the effects of angular momentum quantization, it has been shown (Golubkov {\\it et al.} \\cite{Golubkov}) that the meson emission is modified. Those ideas have not yet been applied to baryonic evaporation. Several improvements could be expected for the detection of antimatter from {\\sc pbh}s in the years to come. First, the AMS experiment (Barrau \\cite{Barrau3}) will allow, between 2005 and 2008, an extremely precise measurement of the antiproton spectrum. In the meanwhile, the BESS (Orito {\\it al.} \\cite{Orito2}) and PAMELA (Straulino {\\it et al.} \\cite{PAMELA}) experiments should have gathered new high-quality data. The solar modulation effect should be taken into account more precisely to discriminate between primary and secondary antiprotons as the shape of each component will not be affected in the same way. The effect of polarity should also be included (Asaoka {\\it et al.} \\cite{Asaoka}). To conclude, the antideuteron signal should also be studied as it could be the key-point to distinguish between {\\sc pbh}-induced and {\\sc susy}-induced antimatter in cosmic rays. Although the $\\bar{p}$ emission due to the annihilation of supersymetric dark matter would have nearly the same spectral characteristics than the {\\sc pbh} evaporation signal, the antideuteron production should be very different as coalescence schemes usually considered (Chardonnet {\\it et al.} \\cite{Chardonnet}) cannot take place between successive {\\sc pbh} jets." + }, + "0112/astro-ph0112165_arXiv.txt": { + "abstract": "We report the first X-ray detection of the radio pulsar PSR B1757$-$24 using the {\\it Chandra X-ray Observatory}. The image reveals point-source emission at the pulsar position, consistent with being magnetospheric emission from the pulsar. In addition, we detect a faint tail extending nearly 20$''$ east of the pulsar, in the same direction and with comparable morphology to the pulsar's well-studied radio tail. The X-ray tail is unlikely to be emission left behind following the passage of the pulsar, but rather is probably from synchrotron-emitting pulsar wind particles having flow velocity $\\sim$7000~km~s$^{-1}$. Assuming the point-source X-ray emission is magnetospheric, the observed X-ray tail represents only $\\sim$0.01\\% of the pulsar's spin-down luminosity, significantly lower than the analogous efficiencies of most known X-ray nebulae surrounding rotation-powered pulsars. ", + "introduction": "PSR~B1757$-$24 is a 124-ms radio pulsar near the supernova remnant (SNR) G5.4$-$1.2. Radio timing observations show that the pulsar has characteristic age 16~kyr and spin-down luminosity $\\dot{E} = 2.6 \\times 10^{36}$~erg~s$^{-1}$ (Manchester et al. 1991). The pulsar is at the tip of a flat-spectrum radio protuberance just outside the west side of the SNR shell. The protuberance consists of a small, roughly spherical nebula, G5.27$-$0.90, that has a highly collimated finger of emission on its western side. The pulsar is at the westernmost tip of the finger, whose morphology and spectrum suggest a ram-pressure confined pulsar wind nebula (PWN) (Frail \\& Kulkarni 1991). Assuming the pulsar was born at the center of G5.4$-$1.2 and that the characteristic pulsar age is a good estimate for the age of a system, the pulsar appears to have overtaken the expanding shell, implying a transverse space velocity of $v_t \\sim$1800~km~s$^{-1}$ for a distance of 5~kpc (Frail, Kassim, \\& Weiler 1994). However, recent interferometric observations failed to detect the implied proper motion (Gaensler \\& Frail 2000). They set a 5$\\sigma$ upper limit of $v_t < 590$~km~s$^{-1}$. This suggests that the pulsar is older than its characteristic age, or that the assumed pulsar birth place is incorrect. A ram-pressure confined wind should radiate X-rays as part of the broad-band synchrotron spectrum that results from the shock-acceleration and subsequent gyration of relativistic wind electron/positron pairs in the ambient magnetic field. This is in contrast to static PWNe in which a high external pressure, usually supplied by the hot gas interior of a SNR, does the confining. The PSR~B1757$-$24 PWN is of interest as it exhibits the cleanest and most extreme bow-shock plus tail morphology of any such system, suggesting that it may exemplify the class of ram-pressure confined PWN most ideally. We summarize here the first X-ray detection of PSR~B1757$-$24. Details of this result can be found in Kaspi et al. (2001). ", + "conclusions": "The X-ray point source is likely to be non-thermal pulse-phase-averaged magnetospheric emission from the radio pulsar itself. The observed 2--10~keV unabsorbed luminosity, for a distance $d=5$~kpc is $2 \\times 10^{33}$~erg~s$^{-1}$, assuming beaming angle $\\phi = \\pi$~sr (see Kaspi et al. 2001). This implies an efficiency of conversion of spin-down luminosity into magnetospheric emission of $0.00020(\\phi/\\pi \\; {\\rm sr})(d/5 \\; {\\rm kpc})^2$. This efficiency, as well as the measured power-law photon index, are consistent with those observed for the magnetospheric components of other radio pulsars (Becker \\& Tr\\\"umper 1997). The tail X-ray emission is likely to be synchrotron radiation from the shocked pulsar wind. The observed X-ray tail extends nearly 20$''$, or 0.48~pc for a distance of 5~kpc, to the east of the pulsar, nearly as long as the detected radio tail. The upper limit on the transverse velocity of $v_t <590$~km~s$^{-1}$ implies the time since the pulsar was at the eastern-most tip of the observed X-ray emission must be $> 800$~yr. The synchrotron lifetime of a photon of energy $E$ (in keV) in a magnetic field $B_{-4}$ (in units of $10^{-4}$~G) is $t_s \\simeq 40 E^{-1/2} B_{-4}^{-3/2}$~${\\rm yr}$. Thus, for $t_s > 800$~yr and $E\\simeq 1-9$~keV, $B< 0.8-14 \\; \\mu$G. This is much less than the equipartition magnetic field $B_{\\rm eq} \\sim \\sqrt{\\dot{E}/r_s^2 c } \\sim 70 \\; \\mu$G expected in the vicinity of the pulsar. Here, $r_s$ is the distance from the pulsar to the bow shock head. Hence, the X-ray tail behind PSR~B1757$-$24 cannot be synchrotron emission from pulsar wind particles just left behind after the passage of the pulsar. Rather, freshly shocked wind particles must be continuously fed eastward with a velocity much larger than the pulsar space velocity, $v_f \\gg v_t$. We can constrain the flow velocity $v_f$ of the wind particles in the tail by noting that it must be high enough to continuously supply particles given their cooling times. Thus, the flow time $t_f \\lapp t_s$. Assuming that the magnetic field reaches its equipartition value near the bow-shock head and that this value holds for the approximately one-dimensional tail region too, we find $t_s \\gapp 70 (E / 1 \\; {\\rm keV})^{-1/2} (B_{\\rm eq} / 70 \\; \\mu{\\rm G})^{-3/ 2}$~yr. As the tail extends to 0.48~pc for $d=5$~kpc, this implies $v_f \\gapp 6700 (d/5 \\; {\\rm kpc})(E / 1 \\; {\\rm keV})^{1/2} (B_{\\rm eq} / 70 \\; \\mu{\\rm G})^{3/2}$~km~s$^{-1}$. The tail emission has low flux. The tail surface brightness in the 2--8~keV band is $4.5 \\times 10^{-16}$~erg~s$^{-1}$~cm$^{-2}$~arcsec$^{-2}$, with uncertainty of $\\sim$30\\%. The total unabsorbed flux in the 2--8~keV band in our extraction region is $9.8 \\times 10^{-14}$~erg~s$^{-1}$~cm$^{-2}$, with similar uncertainty. The efficiency with which the pulsar's $\\dot{E}$ is converted into tail X-rays in the 2--8~keV band is only 0.00011$(d/5 \\; {\\rm kpc})^2$, roughly half of the point-source efficiency. This is in contrast to other rotation-powered pulsars, like the Crab, whose X-ray nebular emission is much brighter than the point-source output. Without timing information, we cannot rule out an ultra-compact nebula as the source of the point-source emission. Several effects may reduce the X-ray efficiency of ram pressure confined PWNs. First, the efficiency of conversion of $\\dot{E}$ into X-ray emitting particles may be lower since the reverse shock in the ram-pressure-confined PWNs is strong only in the forward part of the head, subtending a much smaller solid angle than in a static PWN. Second, a low X-ray efficiency is expected if the flow time of the relativistic plasma through the tail is shorter than the synchrotron life time. A similar argument was put forth (Chevalier 2000) to explain the low efficiencies of the Vela and CTB~80 pulsars. Finally, low surface brightness emission from beyond the eastern tip of the observable X-ray tail, or even from G5.27$-$0.90, have gone undetected in our observation. Emission from the direction of that nebula, having X-ray surface brightness half of the ACIS-S3 background, would contribute roughly two orders of magnitude more flux. No emission is detected from the shell supernova remnant G5.4$-$1.2. The upper limits on remnant emission are unconstraining." + }, + "0112/astro-ph0112414_arXiv.txt": { + "abstract": "We have performed a search for halo white dwarfs as high proper motion objects in a second epoch Wide Field Planetary Camera 2 image of the Groth-Westphal strip. The survey covers 74.8 arcmin$^2$, and is complete to $V\\sim26.5$. We identify 24 high proper motion objects with $\\mu > 0.014^{\\prime\\prime}$/y. Five of these high proper motion objects are identified as strong white dwarf candidates on the basis of their position in a reduced proper motion diagram. We also identify two marginal candidates whose photometric errors place them within $\\sim1\\sigma$ of the white dwarf region of the reduced proper motion diagram. We create a model of the Milky Way thin disk, thick disk and stellar halo and find that this sample of white dwarfs is clearly an excess above the $\\leq 2$ detections expected from these known stellar populations. The origin of the excess signal is less clear. Possibly, the excess cannot be explained without invoking a fourth galactic component: a white dwarf dark halo. Previous work of this nature has separated white dwarf samples into various galactic components based on kinematics; distances, and thus velocities, are unavailable for a sample this faint. Therefore, we present a statistical separation of our sample into the four components and estimate the corresponding local white dwarf densities using only the directly observable variables, $V$, $(V-I)$, and $\\vec{\\mu}$. For all Galactic models explored, our five white dwarf sample separates into about 3 disk white dwarfs and 2 halo white dwarfs. However, the further subdivision into the thin and thick disk and the stellar and dark halo, and the subsequent calculation of the local densities are sensitive to the input parameters of our model for each Galactic component. Using the lowest mean mass model for the dark halo and the 5 white dwarf sample we find $\\nthin = 2.4^{+0.7}_{-0.6}\\times10^{-2}$ pc$^{-3}$, $\\nthick = 0.0^{+7.6}\\times10^{-4}$ pc$^{-3}$, $\\nstellar = 0.0^{+7.7}\\times10^{-5}$ pc$^{-3}$, and $\\ndark = 1.0^{+0.4}_{-0.4}\\times10^{-3}$ pc$^{-3}$. This implies a 7\\% white dwarf halo and six times the canonical value for the thin disk white dwarf density (at marginal statistical significance), but possible systematic errors due to uncertainty in the model parameters likely dominate these statistical error bars. The white dwarf halo can be reduced to $\\sim1.5$\\% of the halo dark matter by changing the initial mass function slightly. The local thin disk white dwarf density in our solution can be made consistent with the canonical value by assuming a larger thin disk scaleheight of 500 pc. ", + "introduction": "The microlensing results towards the Large Magellanic Cloud \\citep{alc97, alc00, lass00} generated much interest in the possibility of white dwarfs (WD) as significant contributors to the Galactic halo dark matter. The most recent results suggest a most likely MACHO fraction of 20\\% and a most likely MACHO mass between 0.15 and 0.9 $M_\\odot$ \\citep{alc00}. Other MACHO candidates such as brown dwarfs, M dwarfs, and neutron stars are excluded, respectively, on the basis of the most likely MACHO mass, direct star counts which suggest that their contribution to the total mass is insignificant \\citep{gou98} and entirely unacceptable nucleosynthesis yields from their main sequence precursors \\citep{car94}. White dwarfs are well-known low luminosity stars, and have been extensively surveyed, for instance by \\citet{leg98} and \\citet{kno99}. A very recent report \\citet{maj01} suggests the scale height for thin disk white dwarfs may be higher than previously thought, and that consequently the total number of thin disk white dwarfs is also higher than previously thought. White dwarfs as dark matter pose their own set of problems, as their main sequence progenitor may produce more metals (He, C, N) than observed \\citep{fie00}. These chemical evolution constraints can possibly be avoided by assuming a non-standard initial mass function \\citep{cha96,cha99} and, perhaps, lower metal yields from Z=0.0 zero age main sequence progenitors. Recent results \\citep{mar01} suggest that the first-dredge up does not take place for zero metallicity stars with $M\\gtrsim1.2~M_{\\odot}$. Further, the second dredge-up is suppressed for zero metallicity stars with $M\\lesssim2.1~M_{\\odot}$ and only brings CNO to the surface for $2.7~M_{\\odot}\\lesssim M \\lesssim 8.3~M_{\\odot}$. Finally, thermal pulses on the asymptotic giant branch which would normally bring carbon to the surface may not occur \\citep{cha99,mar01}. Interest in WD as directly detectable dark matter intensified with the suggestion that ancient hydrogen atmosphere WD evolve towards \\textit{bluer} optical colors as they cool \\citep{han98}, remaining detectable in the $V$ band for many Gyr longer than previously assumed. This provided a plausible explanation for some of the faint blue objects in the Hubble Deep Field (HDF), two of which were reported to have proper motions consistent with an interpretation as ancient halo white dwarfs \\citep{iba99}. Although the HDF moving objects were determined to be false detections \\citep{ric01}, excitement over the possibility of halo white dwarfs was renewed with the results of \\citet{opp01} who claim detection of 38 halo white dwarfs, constituting at least 2\\% of the Galactic halo dark matter. \\citet{opp01} present their results in the form of a plot of the galactic radial (U) and rotational (V) velocities of each WD and superimpose 1 and 2 sigma contours for the expected locations of the thick disk and halo components of the Milky Way. White dwarfs lying outside the 2 sigma contours of the thick disk are assumed to belong to a halo population. \\citet{rei01} provide an alternate interpretation of the \\citet{opp01} results, arguing the velocity distribution is more consistent with the high-velocity tail of the thick disk. Furthermore, by comparing the placement of the halo candidates along a fiducial WD evolution track in a color-magnitude diagram, \\citet{han01} notes that the \\citet{opp01} halo population seems to have an age distribution similar to the standard thin disk population. \\citet{han01} finds that it is difficult to make the age distribution of this sample consistent with common assumptions about thick disk star formation (a simple burst at early times), and even more difficult to achieve consistency with some sort of truly ancient halo population. \\citet{koo01} extend this argument, noting that the thick disk and halo WD populations (as divided by \\citet{opp01}) are indistinguishable in terms of luminosity, color and apparent age. \\citet{koo01} undertake a more sophisticated analysis of the \\citet{opp01} sample, calculating the contribution of the thick disk and halo using a maximum likelihood analysis. They find a local number density of thick disk WD of $n_{0,{\\rm thick~disk}}= 1.8\\pm0.5 \\times10^{-3}$ pc$^{-3}$ and a local number density of halo WD of $n_{0,{\\rm halo}}= 1.1^{+2.1}_{-0.7}\\times10^{-4}$ pc$^{-3}$. The halo density is about 5 times higher then previously expected \\citep{gou98}, but constitutes only $\\sim$ 0.8\\% of the dark halo density, at least an order of magnitude smaller than the MACHO density implied by the microlensing results. \\citet{rey01} also provide a reanalysis of the \\citet{opp01} sample. These authors use a comprehensive description of Galactic stellar populations to provide a simulation of the \\citet{opp01} data, including the detection limits in proper motion, magnitude and color. They conclude that thick disk white dwarfs with standard local densities are sufficient to explain the \\citet{opp01} sample. The wide spread in age of the \\citet{opp01} sample has led several parties to suggest that these WD originated in the thin disk, but were subsequently accelerated to much higher velocities. \\citet{koo01} suggest a mechanism to eject WD into the halo with the required speeds of $\\sim200$ km/s through the orbital instability of triple systems. \\citet{dav01} proposes another binary driven mechanism in which the secondary of a tight binary system (the WD progenitor) is ejected at high velocity when the primary explodes as a supernova. The \\citet{opp01} results are based on a very wide survey (10\\% of the sky) with a bright limiting magnitude ($R_{\\rm lim} \\lesssim 20$). In this work we take the opposite approach, examining a very small area of the sky (20 Wide Field Planetary Camera 2 fields), but probing to a very deep limiting magnitude ($V_{\\rm lim} \\sim 27$). Based on the \\citet{opp01} results which are $\\sim90$\\% proper motion limited \\citep{koo01}, we would not expect the white dwarf luminosity function to rise suddenly beyond their detection limit. However, the \\citet{opp01} survey cannot exclude the existence of an ancient ``dark halo'' white dwarf population of age $>> 10$ Gyr. In fact, since \\citet{opp01} use the \\citet{ber97} survey of young white dwarfs to derive a linear relationship between absolute magnitude and color, they assume that all detected white dwarfs are younger than the beginning of the cooling turnoff towards bluer colors in the color magnitude diagram. In their filters, this color turnoff occurs at a temperature of around 2500 K, or a white dwarf age of about 13-14 Gyr. Since they observe very few white dwarfs near the color turnoff this is probably a good, if limiting, assumption. Although we initially explore an intuitive analysis of our sample along the same lines as \\citet{opp01}, in our final analysis we avoid making any assumptions about the age of our sample and estimate the local densities of disk and halo white dwarfs based solely on the observed properties of our sample: apparent magnitude, color and proper motion vector on the sky. This work is structured as follows. In \\S2 we describe the observation and reduction of each epoch, our procedure for matching stars between epochs, the selection of significant high proper motion objects and the results of our completeness tests. In \\S3 we describe our selection criteria for white dwarfs and examine their kinematic properties in a manner similar to \\citet{opp01}. In \\S4 we model various components of the Milky Way and compute the number of white dwarf detections we expect in our survey. Then, using only the directly observable properties, we create a method to statistically separate our sample into the various Milky Way components. We conclude in \\S5. ", + "conclusions": "\\citet{koo01} and \\citet{rei01} demonstrate that the \\citet{opp01} results do not necessarily imply the presence of a substantial white dwarf dark matter halo. They also note that due to the color turnoff of ancient white dwarfs a survey with such a bright limiting magnitude cannot exclude the presence of an ancient white dwarf halo of age $\\gg10$ Gyr. To exclude an ancient halo requires a survey with a substantially fainter limiting magnitude, such as this one. For faint or ancient white dwarfs, spectroscopic velocities are not available, parallax measurements cannot be made and photometric distance relations calibrated at bright magnitudes are potentially unreliable. Therefore, the division of a faint white dwarf sample into disk and halo contributions must be done without true kinematic (velocity) information. In this work, we have explored such a statistical separation using only the directly observable quantities of the five high proper motion white dwarf candidates detected in a second WFPC2 epoch of the Groth-Westphal strip. The small sample size and our imperfect knowledge of the characteristics of the putative white dwarf dark halo lead to possible large systematic errors in our analysis. Using the 96IMF1 dark halo gives a 7\\% white dwarf halo and $\\nthin= 2.4^{+0.5}_{-0.4}\\times10^{-2}$ pc$^{-2}$, a gross excess of thin disk white dwarfs. However, we explore several alternative Galactic models which demonstrate that uncertainties in our models lead to possible systematic errors which may be larger then the quoted statistical errors. For instance, the thin disk local density can be lowered four-fold and brought to within $1$ sigma of the canonical value by assuming the larger thin disk scale height suggested by \\citet{maj01}. Also, if we assume a slightly different dark halo initial mass function, the halo signal shifts to the stellar halo where it implies a local stellar halo white dwarf density of $\\nstellar \\sim 2.2 \\times 10^{-4}$, about 10 times higher then the canonical value and similar to the results reported by \\citet{opp01}. The use of external constraints also affects our results. If we constrain the thin disk contribution to the canonical value, or attempt to remove the thin disk contribution by hand (as in \\citet{koo01}), the disk signal shifts into the thick disk, giving $\\nthick \\sim 3 \\times 10^{-3}$ pc$^{-3}$. This elevated thick disk density is similar to the results of \\citet{koo01}. However, regardless of the details of our models, we always find $\\sim3$ disk white dwarfs and $\\sim2$ halo white dwarfs, a clear excess above the $\\sim$1--2 total detections expected from known stellar populations. We are unable to definitely determine the source of the excess signal, but its existence seems clear. We summarize the local densities found in our various calculations in Table 3. We note that the ages, abundances and kinematics of the known stellar populations of our Galaxy are still a source of lively debate. We have attempted to highlight how our results depend on our assumptions about these quantities in \\S4.4. Attempts to more precisely determine the properties of the various components of our Galaxy will benefit from surveys with very precise astrometry. Both our results and those of \\citet{opp01} are intriguing and together explore the two extremes of survey philosophies: bright limiting magnitude with many objects and faint limiting magnitude with few objects. The \\citet{opp01} sample is bright enough to be able to employ photometric distance calibrations, however, fainter surveys necessary to truly exclude or confirm the dark white dwarf halo must employ an approach more similar to ours. Future improved knowledge of the input model characteristics and a faint sample with more objects may reduce the model dependencies which hamper our conclusions." + }, + "0112/astro-ph0112308_arXiv.txt": { + "abstract": "We present our recently developed 3-dimensional chemodynamical code for galaxy evolution. This code follows the evolution of different galactic components like stars, dark matter and different components of the interstellar medium (ISM), i.e. a diffuse gaseous phase and the molecular clouds. Stars and dark matter are treated as collisionless N-body systems. The ISM is numerically described by a smoothed particle hydrodynamics (SPH) approach for the diffuse gas and a sticky particle scheme for the molecular clouds. Additionally, the galactic components are coupled by several phase transitions like star formation, stellar death or condensation and evaporation processes within the ISM. As an example we show the dynamical and chemical evolution of a star forming dwarf galaxy with a total baryonic mass of $2 \\cdot 10^9 M_\\odot$. After a moderate collapse phase the stars and the molecular clouds follow an exponential radial distribution, whereas the diffuse gas shows a central depression as a result of stellar feedback. The metallicities of the galactic components behave quite differently with respect to their temporal evolution as well as their radial distribution. Especially, the ISM is at no stage well mixed. ", + "introduction": " ", + "conclusions": "" + }, + "0112/astro-ph0112552_arXiv.txt": { + "abstract": "DEEP\\index{DEEP} is a two-phase spectral survey of faint field galaxies with the Keck Telescopes. The goals include exploring galaxy formation \\index{galaxy formation} and evolution \\index{galaxy evolution}, mapping distant large scale structures, and constraining cosmology. DEEP, since its inception in the early 1990's, has been distinguished by an emphasis on studying the kinematics \\index{kinematics} and masses \\index{masses} of distant galaxies. The major DEEP survey in the second phase (DEEP2) is scheduled to begin in 2002 and will mainly aim for a sample of 50,000 galaxies to $I \\sim 23$. Until then, the first phase of DEEP science programs will have been concentrating on using existing Keck spectrographs to undertake spectral surveys of over 1000 galaxies that have also been observed with HST. I will highlight the study of rotation curves \\index{rotation curves} of distant spirals; the fundamental plane \\index{fundamental plane} of faint, high-redshift E/S0s\\index{ellipticals} \\index{spheroids} \\index{early-type galaxies}; the narrow velocity widths seen in luminous blue compact galaxies; and the diversity of kinematics seen in a small sample of high redshift ($z \\sim 3$) galaxies. These DEEP pilot programs have clearly demonstrated the feasibility, importance, and potential of using kinematics to better understand distant galaxies. ", + "introduction": "The first decade of the 21st century promises many new surveys of distant galaxies, especially with the advent of a suite of new 8-10~m class, ground-based, optical telescopes. Besides adding critical redshifts to data from space and other wavebands, the higher S/N and spectral-resolutions affordable with 8-10m telescopes provide three new and quite powerful diagnostics for the analysis of distant galaxies: internal velocities (i.e., kinematics and hence dynamical masses when size is also measured); chemical abundances; and star-formation-rate and stellar-population-age estimates. Compared to the traditional parameters of counts, colors, luminosities, and clustering properties of distant galaxies, these new diagnostics yield independent probes of galaxy properties in the early universe and have solid links to theoretical simulations of galaxy formation. Moreover, since both galaxy evolution and their large scale patterns involve a complex interplay of diverse galaxy classes, environments, and physical mechanisms and because precision cosmology via the volume test requires averaging over the fluctuations due to large-scale clustering, very large samples are essential to extract reliable results. To meet the challenge, DEEP\\footnote{DEEP: Deep Extragalactic Evolutionary Probe; more details on participants and programs of DEEP can be found at URL: {\\bf http://www.ucolick.org/ $\\sim$deep}} was initiated over 9 years ago as a spectral survey of 50,000 \\footnote{our original goal of 10,000 has been revised upwards to improve significantly the reliability of cosmological tests and of large scale structure studies} faint field galaxies , using the Keck II 10-m Telescope with a new spectrograph DEIMOS \\cite{Dav98} \\footnote{ DEIMOS: DEep Imaging Multi-Object Spectrograph; more information is provided at URL: {\\bf http://www.ucolick.org/ $\\sim$loen/Deimos/deimos.html}}. The use of DEIMOS provides a clean division of DEEP into two parts or phases. The first is a set of pilot-style surveys of relatively small samples (10's - 1000) of galaxies. These pilot surveys exploit the pre-DEIMOS spectrographs available on Keck and were designed to determine feasibility and to refine the scope of DEEP2. DEEP is distinguished by aiming to gather internal kinematic data in the form of rotation curves or linewidths, as well as spectral-line measurements sensitive to star formation rates, gas conditions, stellar-population ages, and metallicity. ", + "conclusions": "The main theme that arises from our DEEP pilot programs is that galaxy evolution is a complex problem. Galaxies are diverse in size, luminosity, structure, etc.; are composed of subcomponents which may experience different star formation and dynamical histories and evolution; and reside in a wide range of environments involving different physical mechanisms for their evolution. We have established that kinematics are both feasible with 8-10~m class telescopes and valuable for understanding distant galaxies. For example, we find relatively little evolution in the Tully-Fisher relation or disk surface brightness \\cite{Sim99} to redshifts $z \\sim 1$, as well as little evidence for evolution in the fundamental plane, volume density, or luminosity beyond that expected from passive evolution for early-type galaxies to $z \\sim 1$. The colors of the spheroids and bulges are, however, redder than expected and thus a puzzle. On the other hand, luminous blue compact galaxies appear, whether at low redshifts $z < 1$ or at high redshifts $z \\sim 3$, to have very low dynamical masses and are suggested to be possible progenitors of quiescent low-mass spheroidals today or the building blocks of larger galaxies rather than massive ellipticals undergoing formation via monolithic collapse. The lessons from our DEEP phase-one pilot programs indicate great promise for our main survey DEEP2 of 50,000 galaxies. Such large numbers are vital for analysis after subdivision of the full sample by a wide range in luminosity, size, M/L, structure, redshift, and environment. More relevant to this conference, we are optimistic that the kinematic data will yield new studies that rely on mass, including mass functions; M/L functions; Tully-Fisher, Fundamental Plane, and other scaling law evolution; mergers rates; dark matter distributions (halo vs disk; large scale structure vs mass); and precision cosmology (velocity function vs volume tests; estimates of the equation of state \\cite{New00}). {\\bf Acknowledgements:} DEEP was initiated by the Berkeley Center for Particle Astrophysics (CfPA), and has been supported by various other NSF, NASA, UC, and STScI grants. The senior members of DEEP have managed the project, but I would like to give special thanks to our talented pool of more junior astronomers over the years (see DEEP URL for names), without whom the results presented here would not have been possible." + }, + "0112/astro-ph0112078_arXiv.txt": { + "abstract": "Most formation scenarios of globular clusters assume a molecular cloud as the progenitor of the stellar system. However, it is still unclear, how this cloud is transformed into a star cluster, i.e.\\ how the destructive processes related to gas removal or low star formation efficiency can be avoided. Here a scheme of supernova (SN) induced cluster formation is studied. According to this scenario an expanding SN shell accumulates the mass of the cloud. This is accompanied by fragmentation resulting in star formation in the shell. Provided the stellar shell expands sufficiently slow, its self-gravity stops the expansion and the shell recollapses, by this forming a stellar system. I present N-body simulations of collapsing shells which move in a galactic potential on circular and elliptic orbits. It is shown that typical shells ($10^5$ M$_\\odot$, 30 pc) evolve to twin clusters over a large range of galactocentric distances. Outside this range single stellar systems are formed, whereas at small galactocentric distances the shells are tidally disrupted. In that case many small fragments formed during the collapse survive as single bound entities. About 1/3 of the twin cluster systems formed on circular orbits merge within 400 Myr. On elliptic orbits the merger rate reduces to less than 4\\%. Thus, there could be a significant number of twin clusters even in our Galaxy, which, however, might be undetected as twins due to a large phase shift on their common orbit. ", + "introduction": "The exact formation process of globular clusters is still under debate. Suggested mechanisms include -- among other scenarios -- e.g.\\ the collapse of giant molecular clouds (GMC) or the collision of molecular clouds (e.g.\\ \\inlinecite{fall85}, \\inlinecite{murray90}, \\inlinecite{fujimoto97}). A common feature of most scenarios is the assumption of smooth initial gas distributions which are transformed into the cluster. However, this assumption requires short formation timescales and unusually high star formation efficiencies in order to end up with a gravitationally bound system. An alternative model introduced by \\inlinecite{brown91} can overcome these difficulties: their scenario starts with an OB-association exploding near the center of a molecular cloud. The expanding shell sweeps up the cloud material and in a later stage the expansion is decelerated and stopped by the accumulated mass as well as the external pressure of the ambient interstellar medium. The shell itself is assumed to undergo fragmentation and, finally, star formation. If these stars form a gravitationally bound system, this stellar shell will recollapse, by this creating a star cluster. At the moment a discrimination between different scenarios by direct simulations (starting from first principles) is far out of reach. However, one can study different evolutionary stages in some detail. E.g.\\ \\inlinecite{theis00} compared in a series of N-body simulations the collapse of thin stellar shells and homogeneous spheres in a galactic tidal field. These calculations were performed for circular and eccentric orbits, but with a constant apogalacticon of 5 kpc. It was found that collapsing shells preferably end in multiple systems, mainly twins, whereas homogeneous spheres either form single clusters or become completely disrupted. In this paper the influence of the galactocentric distance, i.e.\\ the strength of the tidal field, on the collapse of stellar shells is investigated. Special forcus is put to the survival probability of the formed multiple stellar cluster systems. \\begin{figure}[t] \\centerline{\\hbox{ \\psfig{figure=theisc_fig1.eps,width=10.0cm,angle=0} }} \\caption{Snapshots at $t=4\\approx 30 \\,\\, \\mbox{\\rm Myr}$ for circular orbits at different galactocentric distances: $R=2 \\,\\, \\mbox{\\rm kpc}$ (upper left), $R=3.5 \\,\\, \\mbox{\\rm kpc}$ (upper right), $R=5 \\,\\, \\mbox{\\rm kpc}$ (lower left), $R=10 \\,\\, \\mbox{\\rm kpc}$ (lower right). The spatial unit is 30 pc. } \\end{figure} ", + "conclusions": "" + }, + "0112/astro-ph0112287_arXiv.txt": { + "abstract": "We present results from two \\chandra/ACIS observations of the so-called Vela ``Bullet D'' region on the eastern limb of the Vela supernova remnant. The Bullet D region is a bright X-ray feature, identified by Aschenbach et al. (1995) from the ROSAT All-Sky Survey, which protrudes beyond the blast wave on the eastern side of the remnant. It has been suggested that this feature is a fragment of supernova ejecta which is just now pushing beyond the position of the main blast wave. An alternate explanation is that the feature is a ``break-out'' of the shock in which inhomogeneities in the ambient medium cause the shock to be non-spherical. The \\chandra\\/ image shows a fragmented, filamentary morphology within this region. The \\chandra\\/ spectra show strong emission lines of O, Ne, and Mg. Equilibrium ionization models indicate that the O and Ne abundances are significantly enhanced compared to solar values. However, non-equilibrium ionization models can fit the data with solar O abundances and Ne abundances enhanced by only a factor of two. The \\chandra\\/ data are more consistent with the shock breakout hypothesis, although they cannot exclude the fragment of ejecta hypothesis. ", + "introduction": "The Vela supernova remnant (SNR) is a large (diameter $\\sim 8$ degrees), nearby supernova remnant associated with the Vela pulsar and is one of the brightest objects in the X-ray sky. Recent work indicates the distance to Vela is only $250\\pm30$~pc (Cha~\\etal\\/ 1999 and Jenkins \\& Wallerstein 1995), making it an ideal candidate for resolving fine structure in X-rays with the {\\em Chandra X-ray Observatory}. \\rosat\\/ All-sky Survey observations showed a complicated morphology and revealed the outer extent of the remnant for the first time. Aschenbach~\\etal\\/~(1995) identified several features protruding beyond what is believed to be the primary blastwave as ``explosion fragments''. The brightest of these features is the so-called ``Bullet D''. Tsunemi~\\etal\\/~1999 and Miyata~\\etal\\/~2001 observed ``Bullet A'' with \\asca\\/ and \\chandra\\/ respectively and detected strong Si emission lines which they concluded was evidence that Bullet~A was indeed composed of ejecta from the original explosion. Moriguchi~\\etal (2001) suggested that the bullet features were more likely the result of the interaction of the SNR shock with an inhomogeneous medium as indicated by the numerous molecular clouds identified by a CO survey of the region. Redman~\\etal\\/ (2000) noted the coincidence of the bright optical filamentary nebula RCW~37 and Bullet~D. We proposed two \\chandra\\/ observations, one at the head of the bullet and the other in the ``wake'' in order to examine the proposed explanations for Bullet~D. \\vspace{-0.15in} ", + "conclusions": "The morphology of the X-ray emission revealed by \\chandra\\/ is more consistent with a SNR shock interacting with an inhomogeneous medium than with a discrete fragment of ejecta. The \\chandra\\/ spectra indicate that Ne is enhanced above solar values, but at a rather modest level, and O is consistent with the solar value. We therefore conclude the \\chandra\\/ observations of the Vela Bullet D region are more consistent with a shock breakout hypothesis than with a bullet of ejecta hypothesis, although the \\chandra\\/ data cannot rule out the bullet idea. \\vspace{0.10in}" + }, + "0112/hep-th0112165_arXiv.txt": { + "abstract": "We study the dynamical equations for extra-dimensional dependence of a warp factor and a bulk scalar in 5d brane world scenarios with induced brane metric of constant curvature. These equations are similar to those for the time dependence of the scale factor and a scalar field in 4d cosmology, but with the sign of the scalar field potential reversed. Based on this analogy, we introduce novel methods for studying the warped geometry. We construct the full phase portraits of the warp factor/scalar system for several examples of the bulk potential. This allows us to view the global properties of the warped geometry. For flat branes, the phase portrait is two dimensional. Moving along typical phase trajectories, the warp factor is initially increasing and finally decreasing. All trajectories have timelike gradient-dominated singularities at one or both of their ends, which are reachable in a finite distance and must be screened by the branes. For curved branes, the phase portrait is three dimensional. However, as the warp factor increases the phase trajectories tend towards the two dimensional surface corresponding to flat branes. We discuss this property as a mechanism that may stretch the curved brane to be almost flat, with a small cosmological constant. Finally, we describe the embedding of branes in the 5d bulk using the phase space geometric methods developed here. In this language the boundary conditions at the branes can be described as a 1d curve in the phase space. We discuss the naturalness of tuning the brane potential to stabilize the brane world system. \\bigskip \\noindent CITA-2001-82 ", + "introduction": "One of the most interesting recent directions in high energy physics phenomenology is the development of brane world scenarios in which our 3+1 dimensional spacetime is a 3-brane embedded in a higher dimensional spacetime. In application to the very early universe this leads to brane world cosmology, where the universe we observe is a 3+1 dimensional curved brane embedded in the bulk. One of the issues in brane world scenarios is the warped geometry of the internal space. In addition to the warp factor in the bulk, brane world scenarios often contain bulk scalar fields. Examples include the dilaton in Horava-Witten theory \\cite{HW} (associated with the volume of the compactified 6d Calabi-Yau space) where the 5d effective theory can be obtained \\cite{Lukas}; the Randall-Sundrum model \\cite{RS1} with phenomenological stabilization \\cite{GW} where the choice of the bulk/brane potentials must be consistent with the 5d warp geometry \\cite{Dewolfe,GKL}, the scalar sector of the supergravity realization of the Randall-Sundrum model \\cite{AVP}, bulk supergravity with domain walls \\cite{ST} and others. The 5d bulk scalar plus gravity brane world system is based on the five dimensional Einstein equations with junction conditions at the branes. Here we consider a simpler problem where the 5d space can be split into 4+1 de~Sitter slices \\begin{equation} \\label{warp} ds^2 = dw^2+ A^2(w)ds_4^2 \\, . \\end{equation} The 4d de~Sitter geometry is described by its scalar curvature ${^4}R=12H^2$, where $H=\\text{const}$ is the 4d Hubble parameter. The warp factor $A(w)$ is determined up to boundary conditions by the five dimensional Einstein equations. For the sake of generality we also present corresponding results for the more general case of a $D$ dimensional warped metric with $D-1$ dimensional de~Sitter slices. The limit of vanishing $H$ corresponds to a flat brane, while nonvanishing $H$ corresponds to brane inflation. The brane sets the boundary conditions for the warp factor $A(w)$ and the scalar field $\\phi(w)$. It is known that often the warped geometry (\\ref{warp}) somewhere outside of the brane encounters a spacetime singularity. One way to cure this problem is to invoke a second brane to screen the singularity by making the inner geometry periodic with the inter-brane interval. Two end-of-the-world branes provide orbifold compactification of the inner space. In the Randall-Sundrum model \\cite{RS1} with AdS bulk geometry without scalars the second brane may be removed. The properties of warped geometry with one or two branes were studied in many papers, see e.g.~\\cite{CR,ST,Dewolfe,GKL,FTW,Davis}. The purpose of this paper is to investigate the global properties of 5d warped geometry (\\ref{warp}) for a variety of bulk scalar field potentials $V(\\phi)$, supplemented by boundary conditions at the branes. We will try to understand how typical is the singularity in the warped geometry, how much tuning is required for the brane potentials, and how these depend on the brane curvature $H^2$. Our approach is different from what was used in the earlier literature. The setting of the problem for the geometry (\\ref{warp}) is similar to the investigation of the FRW universe geometry \\begin{equation} \\label{fr} ds^2 = -dt^2+ a^2(t)ds_3^2 \\, . \\end{equation} with the scale factor $a(t)$ and a scalar field with the potential $V(\\phi)$. Powerful method that has been used to investigate this 4d problem is the construction of phase portraits for the dynamic system for variables $\\phi$, $\\dot{\\phi}$, and $\\dot{a}/a$. Using this method it can be shown that for a broad range of potentials $V(\\phi)$ inflation occurs along a separatrix that is a typical intermediate asymptotic for a broad band of phase trajectories \\cite{acad,KLS}. Inspired by this analogy, we adopt the phase portrait approach to studying the warped geometry of the brane world scenario. It turns out that the equations for the system $(A(w), \\phi(w))$ with the potential $V(\\phi)$ are similar to the cosmological equations for $(a(t), \\phi(t))$ but with the sign of the potential reversed. (There are also differences in the numerical coefficients in 4d and 5d.) Flipping the sign of the potential makes a big difference. For example, it alters the geometry of the phase portrait by connecting branches with positive and negative ``Hubble'' parameter ${A' \\over A}$. This connection with 4d cosmology suggests a convergence of this work with recent work on 4d cosmology with negative potentials \\cite{fastroll}, the results of which can be extended to the warped geometry. Our results overlap with \\cite{fastroll} and the connection will be investigated further \\cite{FFKL}. The structure of the paper is the following. In Section 2, we introduce the basic equations for the brane world scenario. In Section 3, we discuss generic properties of the brane world phase space in terms of ${A' \\over A}, \\phi', \\phi$ and classify its critical points. In Section 4, we systematically construct the phase portrait for a 5d space with flat 4d curvature $H=0$ (flat branes) for the simple quadratic potential $V(\\phi)={1 \\over 2}m^2 \\phi^2$. We will see that without branes all trajectories begin and end at naked singularities dominated by the gradient energy $\\phi'^2$ of the scalar field, which corresponds to a ``stiff'' equation of state with anisotropic pressure. We also consider quadratic potentials with positive and negative cosmological constants. In Section 5, we consider exponential potentials $V(\\phi)=V_0 e^{-2\\sqrt{2}\\phi}$. In Section 6, we extend the method of phase portraits to brane world scenarios with curved branes $H \\ne 0$. We shall see that the brane with the larger warp factor will have smaller curvature. We discuss how this effect may be related to the problem of the small cosmological constant on the visible brane. In Section 7, we derive the Hamilton-Jacobi form of the self-consistent Einstein equations for warped geometry with a scalar field, which leads to the SUSY form of an arbitrary positive bulk scalar potential (without any underlying supersymmetry). This correspondence has been previously noted in context of holographic renormalization group flows \\cite{deBoer:1999}. We also address the similarity of the Einstein-Hamilton-Jacobi constraint equation and the well-known gravitational stability form of the potential \\cite{susy1,susy2,ST}. In Section 8, we introduce branes to screen the singularities. We show how the brane boundary conditions can be represented geometrically as a 1d curve in the 3d phase space of the system. It turns out to be convenient to use the EHJ formalism (in many respects similar to using the SUSY form of the potential). From this perspective we will discuss potentials that lead to brane stabilization and the degree of fine-tuning required to achieve them. The paper concludes with the summary of our results. There is also an appendix in which the locations of critical points at infinity are derived for the phase portraits shown here. ", + "conclusions": "We have developed here a method for systematically exploring the properties of different potentials in brane world warped geometry. To construct the phase portrait of the dynamical system of gravity/scalar, one can apply the qualitative theory of differential equations. Solutions of these equations are represented by the trajectories propagating in the phase space. For a single bulk scalar, trajectories are in three dimensional phase space. For the case of the flat branes, all trajectories are located at two dimensional surface, and the phase portrait of the dynamical system can be easily investigated. In general, the phase space trajectories have timelike singularities at one or two of their ends. These singularities are dominated by the scalar field gradient term, and associated with the infinite critical points in the phase space. We describe how to find critical points for an arbitrary potentials. There are, however, examples of the potentials without singularity at one of the end of phase trajectory. In this case, it is possible to construct non-singular warped geometry with a single brane with $Z_2$ symmetry. We also considered the Einstein-Hamilton-Jacobi formulation of the warped geometry with scalar field. Constrain equation relate the arbitrary bulk potential, the Hamiltonian and its $\\phi$-derivative. Surprisingly, the scalar potential taken on the self-consistent solutions acquires SUSY form even without underlying supersymmetry in the theory. We address the issue how this form of the constrain equation for arbitrary ``on-shell'' potential is related to the requirement to the of the SUSY form of the potential for gravitational stability of gravity/scalar system. One can use the phase space with bulk trajectories to study warped geometry between two branes. Junction conditions for each brane generate one dimensional curve in the phase space. Segment of trajectory between two such curves corresponds to the inter-brane warp factor and scalar field. Without tuning the potential, this configuration in general is not located at the two dimensional surface which represents the flat branes, in other words, the solution exist in general for curved branes. However, one can achieve solution with two flat branes by a simple shift of the potentials. This analysis can be easily extended for more realistic case of several bulk scalar degrees of freedom. For instance, for two bulk scalars, phase space is five dimensional and brane junction conditions generate two dimensional surface. Without tuning the brane potentials, in general there is a segment of trajectory which connect both two dimensional surfaces in 5d (except special cases). We leave the phenomenological applications of our methods for construction of braneworld scenario with stabilization and investigation of their stability for future work." + }, + "0112/astro-ph0112528_arXiv.txt": { + "abstract": "Three-dimensional fractal models on grids of $\\sim200^3$ pixels are generated from the inverse Fourier transform of noise with a power law cutoff, exponentiated to give a log normal distribution of density. The fractals are clipped at various intensity levels and the mass and size distribution functions of the clipped peaks and their subpeaks are determined. These distribution functions are analogous to the cloud mass functions determined from maps of the fractal interstellar medium using various thresholds for the definition of a cloud. The model mass functions are found to be power laws with powers ranging from $-1.6$ to $-2.4$ in linear mass intervals as the clipping level increases from $\\sim0.03$ to $\\sim0.3$ of the peak intensity. The low clipping value gives a cloud filling factor of $\\sim10$\\% and should be a good model for molecular cloud surveys. The agreement between the mass spectrum of this model and the observed cloud and clump mass spectra suggests that a pervasively fractal interstellar medium can be interpreted as a cloud/intercloud medium if the peaks of the fractal intensity distribution are taken to be clouds. Their mass function is a power law even though the density distribution function in the gas is a log-normal. This is because the size distribution function of the clipped clouds is a power law, and with clipping, each cloud has about the same average density. A similar result would apply to projected clouds that are clipped fractals, giving nearly constant column densities for power law mass functions. The steepening of the mass function for higher clip values suggests a partial explanation for the steeper slope of the mass functions for star clusters and OB associations, which sample denser regions of interstellar gas. The mass function of the highest peaks is similar to the Salpeter IMF, suggesting again that stellar masses may be determined in part by the geometry of turbulent gas. ", + "introduction": "Interstellar gas appears scale-free when viewed with Fourier transform power spectra (Crovisier \\& Dickey 1983; Green 1993; Lazarian \\& Pogosyan 2000; St\\\"utzki et al. 1998; Stanimirovic et al. 1999; Elmegreen, Kim, \\& Staveley-Smith 2001), delta variance techniques (St\\\"utzki et al. 1998; Zielinsky \\& St\\\"utzki 1999), spectral correlation functions (Rosolowsky et al. 1999), principal component analysis (Heyer \\& Schloerb 1997), perimeter-area measures (Dickman, Horvath, \\& Margulis 1990; Falgarone, Phillips, \\& Walker 1991), box-counting techniques (Westpfahl et al. 1999), and multifractal analysis (Chappell \\& Scalo 2001). The interstellar medium (ISM) looks like a collection of discrete clouds, however, when intensity contours are drawn (Solomon et al. 1987; Loren 1989; St\\\"utzki \\& G\\\"usten 1990; Williams, de Geus, \\& Blitz 1994; Lemme et al. 1995; Kramer et al. 1998; Heyer, Carpenter, \\& Snell 2001) or when spectral absorption lines are fit to Gaussians (Adams 1949; Hobbs 1978; Clark 1965; Radhakrishnan \\& Goss 1972; Spitzer \\& Jenkins 1975). These two interpretations have led to distinct models for the origin of gas structure and star formation. Scale-free models typically involve turbulence and self-gravity (Falgarone \\& Phillips 1990; Scalo 1990; Pfenniger \\& Combes 1994; Lazarian 1995; Elmegreen 1999b; Rosolowsky et al. 1999; MacLow \\& Ossenkopf 2000; Pichardo et al. 2000; Klessen, Heitsch, \\& MacLow 2000; V\\'azquez-Semadeni, Gazol, \\& Scalo 2000; Semelin \\& Combes 2000; Ostriker, Stone, \\& Gammie 2001; Heitsch, Mac Low, \\& Klessen 2001; Toomre \\& Kalnajs 1991; Wada \\& Norman 1999, 2001). Cloudy models involve sticky collisions and star formation triggered by colliding and compressed clouds (Kwan 1979; Hunter et al. 1986; Tan 2000; Scoville, Sanders, \\& Clemens 1986). For general interstellar gas dynamics, the turbulent model (von Weizsacker 1951; Sasao 1973) may be more realistic than the cloudy model for the origin of structure (LaRosa, Shore \\& Magnani 1999; Ballesteros-Paredes, Hartmann, \\& V\\'azquez-Semadeni 1999). Turbulence also gives cloud-like spectral lines through both density structure and velocity crowding (Ballesteros-Paredes, V\\'azquez-Semadeni, \\& Scalo 1999; Lazarian \\& Pogosyan 2000; Pichardo, et al. 2000). For sudden transition fronts like expanding shells, spiral arms, and dust lanes, cloud collisions may be an appropriate way to model the dynamics (e.g., Kenney \\& Lord 1991; Elmegreen 1988). This is because the pre-front clumps made by turbulence at slow ambient speeds are forced to collide together and interact at much faster speeds inside the front. Strongly self-gravitating clouds that are made by turbulence at moderate speeds but then given a chance to cool and settle to high densities should also interact as in the cloudy models. Because of their high densities, these clouds or clumps should move somewhat independently of the surrounding turbulent gas, perhaps lagging behind the expanding flows to make bright rims, or punching through spiral dustlanes to make feathery structures (Seth 2000). The most obvious point of contact between these two views of gas structure is the mass spectrum of the regions that are isolated enough to be defined as clouds. The scale-free nature of the gas shows up as a scale-free mass spectrum for the clouds (Elmegreen \\& Falgarone 1996; St\\\"utzki et al. 1997). This mass spectrum is not useful in the turbulent model because it does not reflect the continuous distribution of matter that is really present. The spectrum is more important for bound star clusters, which are better defined (Elmegreen et al. 2000), and for individual stars, whose masses are probably proportional to the primordial clump masses (Motte, Andr\\'e, \\& Neri 1998; Testi \\& Sargent 1998; Bacmann et al. 2000; Tachihara et al. 2000). Here we model the cloud mass and size spectra using the intensity peaks in a simulated fractal to represent clouds. The results show the anticipated $\\sim M^{-2}dM$ spectrum that comes from simple hierarchical models (Fleck 1996), but for a different reason than what is usually given. In the usual interpretation, the $M^{-2}dM$ spectrum comes from the fact that there is a constant total mass in each logarithmic interval of mass; i.e., each small clump at the bottom is contained inside each large clump at the top. Then $M\\xi(\\log M)d\\log M=$ constant$\\times d\\log M$ becomes $n(M)dM\\equiv\\xi(\\log M)d\\log M= \\xi(\\log M)dM/M\\propto M^{-2}dM$. We show here that the bottom of the hierarchy, where the gas density has local isolated peaks, has about the same spectrum. This is true even though the probability distribution function of density alone is not a power law, but log-normal. We also find that the spectrum flattens to $M^{-1.6}dM$ as the lower limit to the cloud intensity is decreased. This flattening may explain the difference between the mass spectrum of star clusters, which is close to $M^{-2}dM$ (Battinelli et al. 1994; Elmegreen \\& Efremov 1997; Whitmore \\& Schweizer 1995; Zhang \\& Fall 1999) and the mass spectrum for clouds, which typically ranges between $M^{-1.5}$ and $M^{-1.8}$ (e.g., Blitz 1993; Kramer et al. 1998; Heyer, Carpenter, \\& Snell 2001). An even steeper spectrum results from the highest clipping levels modeled here and is close to the Salpeter IMF for stars. Thus the increase in slope of the mass functions from clouds to clusters to stars may be partly the result of different density thresholds in the same fractal gas. ", + "conclusions": "The connection between the cloud model of interstellar structure and the multifractal model is readily understood when clouds are viewed as isolated peaks in the fractal. Here we have shown that the mass spectrum for such fractal clouds is similar to the observed cloud spectrum when clouds are defined or selected to be those regions where the local density exceeds several percent of the peak (resolved) density, or when clouds are defined to be the resolvable peaks plus a proportional amount of gas in the underlying plateaus. Unresolved peaks can have higher densities, as can a fractal model with more cells, but this will not affect the mass spectrum of the resolved objects. Observational cloud surveys tend to span only a factor of 5 to 10 in physical scale (see references for cloud and clump surveys given above). This limited range is entirely a selection effect because unbiased surveys analyzed with power spectra or similar techniques find a much wider range of scales in the interstellar medium, up to a factor greater than $\\sim100$ in the case of whole galaxies. In the cloud surveys, the factor of $\\sim10$ limit arises because smaller regions are unresolved and larger regions are subdivided into separate clouds, the whole structure being hierarchical (Scalo 1985). For many CO surveys, a factor of 10 in scale corresponds to a factor of 10 in average density (Larson 1981), and a factor of 10 in density corresponds to a volume filling factor of 10\\%. This latter result is because for an infinitely self-similar fractal, the filling factor scales inversely with the density (Elmegreen 1999c). In fact, most cloud surveys report a filling factor of about $\\sim10$\\% (e.g., see review in Blitz 1993); this is only an artifact of sampling in the fractal point of view. In the current models, the average filling factors for the clouds in Section \\ref{sect:model1} are $1.2\\times10^{-4}$, $7.0\\times10^{-3}$, and 0.13 for clipping densities of 0.3, 0.1, and 0.03 times the peak. For the clouds in section \\ref{sect:model2}, they are $1.3\\times10^{-4}$, $5.7\\times10^{-3}$, and 0.065. These filling factors vary approximately as the inverse cube or inverse square of the density threshold because the low clipping levels approach the minimum density, and the sprawling cloud boundaries take up excess volume. The models with the lowest clipping levels, 0.03 and $e^{-3}=0.05$, have the about same filling factors as the observed clouds and clumps in surveys. The point here is that the density range that is inadvertently selected for real cloud surveys is also what the fractal models need to give the observed filling factor and mass function. If different observing techniques are employed, giving a wider range of cloud sizes or average densities, for example, thereby increasing the clump filling factor, then we predict that the slope of the resulting mass function will decrease. Similarly, surveys with molecular tracers sensitive to very high densities should be selecting regions closer to the peaks of the turbulent fractal, with lower filling factors, and these surveys should get slightly steeper slopes for the clump mass spectrum. Note that the absolute density of an observed region is not important for the slope of the mass spectrum, only the relative density compared to the peaks and valleys in that region. As the threshold density approaches the minimum density, the mass spectrum flattens because the cloud boundaries spread out and the low density \"intercloud\" medium, which is just the sub-threshold gas in the same fractal, gets included more and more with each cloud. This interpretation is valid for molecular clouds and other gases with a single phase of thermal temperature. Of course, the intercloud means something different when there is a high temperature phase. For example, fractal models of the neutral hydrogen in the Large Magellanic Cloud require a multi-phase structure because of the high density contrasts that are present (Elmegreen, Kim, \\& Staveley-Smith 2001). The size distribution function for our model clouds is also similar to that for interstellar clouds in the best-fit case. The power law nature of the size spectrum, combined with the cutoff in density in the definition of a cloud, ensures that the mass spectrum is also a power law. This is because the average density is always some factor of order unity times the cutoff. Thus the mass spectrum is a power law even though the density distribution function is a log-normal. The mass spectrum for star clusters is steeper than that for clumps, by several tenths (cf. Sect. 1). The fractal model suggests that this steepness results from the denser cloud regions that are sampled by clusters. The typical density of an embedded star cluster corresponds to an H$_2$ density of around $10^5$ cm$^{-3}$ (Lada, Evans \\& Falgarone 1997). This is much closer to the peak density of the turbulent fractal than the density threshold for the cloud itself, so the mass spectrum should be steeper in this model. However, gravity modifies the gas density in the region of a cluster, so the present results are only suggestive. The densest peaks sampled here, with $>30$\\% of the peak density in the cloud, have mass spectra with slopes of $\\sim-2.3$ and $\\sim-2.4$ for models 1 and 2, respectively. These are similar to the slope of the Salpeter stellar IMF. Indeed, mm-continuum sources and stars do sample denser regions than clusters, so the increase in slope is sensible. However, other factors enter into the stellar IMF, such as the relative rate of collapse in regions with different masses and the competition for mass. Thus the importance of the purely fractal result for the IMF cannot be assessed without further modeling (see Elmegreen 1997, 1999a)." + }, + "0112/astro-ph0112234_arXiv.txt": { + "abstract": "We present deep, narrow-band and continuum images of the powerful high-redshift radio galaxy 3C\\hspace*{0.35ex}265 (z=0.811), taken with the TAURUS Tunable Filter on the William Herschel Telescope, together with detailed long-slit spectroscopic observations along the axis defined by the UV/optical emission elongation. The deep images reveal the existence of cones in the ionization structure of 3C\\hspace*{0.35ex}265 within $\\sim$7 arcsec (58 kpc) of the nucleus, where the emission-line structure is not observed to be closely aligned with the radio axis. This indicates that anisotropic illumination from the central active nucleus dominates on a small scale. In contrast, at larger distances ($\\gtrsim$ 10 arcsec; 80 kpc) from the nucleus, low-ionization emission gas is closely aligned with the radio axis, suggesting that jet-cloud interactions may become the dominant mechanism in the line-emitting gas on a larger scale. Moreover, the presence of a high-velocity cloud at 2.5 arcsec from the nucleus, close to the radio axis, indicates that even close to the nucleus ($\\sim$20 kpc) jet-induced shocks have an important kinematic effect. However, spectroscopic analysis of this region reveals that the ionization state of the high-velocity gas is similar to or higher than that of the surroundings, which is opposite to what we would expect for a cloud that has been compressed and accelerated by jet-induced shocks. Our images show that, while on a large scale the low-ionization emission-line structures are aligned with the radio axis, on a smaller scale, where AGN-photoionization dominates, the highest surface-brightness structure is aligned with the closest companion galaxy (misaligned with the radio axis). This suggest that much of the emission-line structure reflects the intrinsic gas distribution, rather than the ionization pattern imprinted by the radio jets or by illumination from the central AGN. Overall, our results underline the need for a variety of mechanisms to explain the properties of the extended emission-line gas in the haloes of radio galaxies. ", + "introduction": "One of the most important issues concerning radio galaxies is the dominant ionization mechanism of the extended line-emitting gas in these sources. These emission-line regions extend large distances from the nucleus, and provide important information about both the origin of the activity and the origin of the extended gas. However, the main physical processes affecting the properties of the extended emission-line regions (EELR) are not fully understood. The two most accepted models are: a) photoionization by the central active galactic nucleus (AGN), and b) shock-ionization by the interactions between the radio- and line-emitting structures, the so-called `jet-cloud interaction' model. (See \\pcite{cliverew2001} for a review.) While most (but not all) low-redshift radio galaxies appear to be consistent with AGN-photoionization, as we move to higher redshifts there is increasing evidence for shocks. Many distant radio galaxies present highly disturbed kinematics, mostly along the radio axis \\cite{mccarthy96}, and highly collimated optical/UV structures aligned along the radio axis \\cite{mccarthy87,best96}. In addition to the general alignments, detailed study of a sample of z$\\sim$1 radio galaxies has revealed a strong evolution of the emission-line properties with radio size \\cite{best2000}: the emission-line regions of small radio sources show disturbed kinematics and emission-line ratios in agreement with shock predictions; while larger radio sources appear more quiescent and present emission-line ratios consistent with AGN-photoionization. However, recent kinematic studies of a sample of z$\\lesssim$0.8 radio galaxies have provided evidence that shocks have an important effect on the emission-line properties in all sources, even those in which the radio structures are on a much larger scale than the emission-line structures \\cite{carmen2001}. It seems clear that a combination of both AGN-photoionization and shock-ionization is required to explain the observed properties in the emission-line regions of radio galaxies over the whole range of redshifts and radio power. However, the balance between these mechanisms is not yet clear. In order to address this issue, we require complete maps of both the kinematics and ionization structure in the extended haloes of radio galaxies, and not only spectroscopic studies along a preferred direction (e.g. radio axis), which offer a limited view of the properties of the emission-line gas in these sources. To this end, in this paper we present deep emission-line imaging of the powerful high-redshift radio galaxy 3C\\hspace*{0.35ex}265, taken with the TAURUS Tunable Filter (TTF), which are aimed at determining the dominant physical mechanisms. While the published ground-based images of high-redshift radio galaxies are relatively shallow, and the higher-resolution Hubble Space Telescope (HST) images are insensitive to low-surface brightness structures, the TTF is sensitive enough to detect faint structures in the galaxy halo. In addition, the TTF also allows a narrow bandpass, which, for our observations, was tuned to observe the high-velocity gas \\cite{tadhunter91}, as well as the unshifted component, therefore permitting a two-dimensional kinematic analysis of the galaxy. To supplement the TTF images, we also present long-slit spectroscopic observations of 3C\\hspace*{0.35ex}265 taken along PA\\hspace*{0.35ex}145$^{\\circ}$, the axis defined by the UV/optical emission elongation. Preliminary results of these observations were presented in \\scite{tadhunter91}. Here we present a more detailed analysis. Throughout this paper a Hubble constant of H$_{0}$ = 50 km s$^{-1}$ Mpc$^{-1}$ and a density parameter of $\\Omega_{0}$~=~1 are assumed, resulting in an angular scale of 8.25 kpc~arcsec$^{-1}$ for 3C\\hspace*{0.35ex}265. ", + "conclusions": "" + }, + "0112/astro-ph0112144_arXiv.txt": { + "abstract": "About 100 steep spectrum radio sources from the RATAN--600 RC catalog were mapped by the VLA and identified with optical objects down to 24$^m$--25$^m$ in the R band using the 6\\,m telescope. An updated list of calibrators with known redshifts of the same class RGs was compiled to evaluate the accuracy of photometric redshifts estimates. BVRI photometry for 60 RC objects was performed with the 6\\,m telescope, and by standard model fitting we have estimated colour redshifts and ages of stellar population of host gE galaxies. The mean redshift of FRII RGs from the RC list happened to be $\\approx$1. Several objects were found in which active star formation began in the first billion years after the Big Bang. ", + "introduction": " ", + "conclusions": "" + }, + "0112/astro-ph0112534.txt": { + "abstract": "We study the mass--to--light ratio of galaxy systems from poor groups to rich clusters, and present for the first time a large database for useful comparisons with theoretical predictions. We extend a previous work, where $B_j$ band luminosities and optical virial masses were analyzed for a sample of 89 clusters. Here we also consider a sample of 52 more clusters, 36 poor clusters, 7 rich groups, and two catalogs, of $\\sim 500$ groups each, recently identified in the Nearby Optical Galaxy sample by using two different algorithms. We obtain the blue luminosity and virial mass for all systems considered. We devote a large effort to establishing the homogeneity of the resulting values, as well as to considering comparable physical regions, i.e. those included within the virial radius. By analyzing a fiducial, combined sample of 294 systems we find that the mass increases faster than the luminosity: the linear fit gives $M\\propto L_{B}^{1.34 \\pm 0.03}$, with a tendency for a steeper increase in the low--mass range. In agreement with the previous work, our present results are superior owing to the much higher statistical significance and the wider dynamical range covered ($\\sim 10^{12}$--$10^{15}~M_{\\odot}$). We present a comparison between our results and the theoretical predictions on the relation between $M/L_B$ and halo mass, obtained by combining cosmological numerical simulations and semianalytic modeling of galaxy formation. % ", + "introduction": "Since the work by Zwicky (1933), it is well known that the luminous matter associated with galaxies in clusters provides only a small part of the total cluster mass. The relative contribution of the dark matter component is usually specified in terms of the mass--to--light ratio, $M/L$, the total amount of mass relative to the total light within a given scale. Pioneering analyses showed that $M/L$ increases from the bright luminous parts of galaxies to cluster scales (Blumenthal et al. 1984). Indeed, models of biased galaxy formation, where galaxies formed only in the highest peaks in the initial fluctuation spectrum, naturally predict an increase of $M/L$ with system mass (e.g., Bardeen et al. 1986; Davis et al. 1985). Several mechanisms whereby the efficiency of galaxy formation is biased towards very high density peaks are possible (e.g., Rees 1985). However, only recently has the combination of cosmological N--body simulations and semianalytic modeling of galaxy formation allowed realistic predictions about the $M/L$ of galaxy systems (Kauffmann et al. 1999; Bahcall et al. 2000; Benson et al. 2000; Somerville et al. 2001). Although differing in details, it has been generally found that $M/L$ increases with mass halo from very poor to rich systems, possibly with a flattening on large scales. As for the observational point of view, the estimate of $M/L$ in galaxy systems is not an easy task. Both mass and luminosity estimates are fraught with several uncertainties. The uncertainties in the luminosity determination are related to corrections for calibration of the photometry (when using inhomogeneous photometric data), background galaxy contamination, and the need to extrapolate the sum of measured luminosities of galaxy members to include faint galaxies and the outer parts of the systems, beyond the region studied (see, e.g., Oemler 1974). Also the estimate of masses is not an easy task, in spite of the various methods which have been applied (e.g., Narayan \\& Bartelmann 1996; Schindler 1996; Mellier 1999; Biviano 2001). Masses of galaxy systems are inferred from either $X$--ray or optical data, under the general hypothesis of dynamical equilibrium. Estimates based on gravitational lensing do not require assumptions about the dynamical status of the system, but a good knowledge of the geometry of the potential well is necessary. Claims for a discrepancy (by a factor of 2--3) between cluster masses obtained with different methods cast doubts about the general reliability of mass estimates (e.g., Wu \\& Fang 1997). However, recent analyses have shown that, if we avoid cases of bimodal clusters, mass estimates concerning large cluster areas are in general agreement (Allen 1997; Girardi et al. 1998b, hereafter G98; Lewis et al. 1999). Large collections of observational data concerning galaxies, groups, and clusters suggest that all systems have a constant ratio of $M/L_B\\sim 200$--$300$ \\ml for scales larger than galaxies, so that the total mass of galaxy systems could be roughly accounted for by the total mass of their member galaxies, possibly plus the mass of the hot intracluster gas (Rubin 1993; Bahcall, Lubin, \\& Dorman 1995). Homogeneous samples, where both masses and luminosities are computed in a consistent way, would be more reliable. Unfortunately, the above observational difficulties prevented us from building a large $M/L$ data base spanning a wide dynamical range. Based on homogeneous optical data, the pioneering work by Dressler (1978) showed no evidence of correlation of $M/L$ values with richness for 12 clusters. More recently, David, Jones, \\& Forman (1995), who used homogeneous X--ray mass estimates and luminosities from different sources in the literature, showed that $M/L_V$ of seven groups and clusters of galaxies are comparable. Also $M/L_r$ values for the sample of 15 clusters of the Canadian Network for Observational Cosmology (CNOC, Calberg et al. 1996), where masses come from optical virial estimates, are consistent with an universal underlying value. A slight increase of $M/L$ with mass system was suggested by indirect analyses of the cluster fundamental plane, i.e., the study of the relations between cluster size, internal velocity dispersion, and luminosity (but see Fritsch \\& Buchert 1999). In fact, assuming the virialization state and internal structure of all clusters to be identical, one can derive the behavior of $M/L$. Working with a homogeneous photometric sample of 12 clusters Schaeffer et al. (1993) found that $M/L_V\\propto L_V^{0.3}$. Similarly, using homogeneous results for 29 clusters of the ESO Nearby Abell Cluster Survey (ENACS; Katgert et al. 1998), Adami et al. (1998a) showed that $M/L_{B_j}\\propto \\sigma_v$, where $\\sigma_v$ is the line--of--sight (l.o.s.) velocity dispersion of member galaxies. This correlation was also directly verified by Adami et al. (1998b) in a following work by computing the projected virial masses. Recently, Girardi et al. (2000, hereafter G00) faced the question from a direct point of view with a significant increment in the data--base statistics. They analyzed 89 clusters, all with homogeneous optical virial mass estimates by G98 and homogeneous luminosity estimates derived from the COSMOS catalog (Yentis et al. 1992). Moreover, the available data allowed the authors to compute mass and luminosity within the virial radius in order to analyze physically comparable regions in poor and rich clusters. Their main result is that the mass has a slight, but significant tendency to increase faster than the luminosity: $M\\propto L_{B_j}^{\\mbox{\\rm 1.2--1.3}}$. Owing to the large uncertainties generally involved, it is really not surprising that such a slight effect could not be detected by previous analyses based on small statistics and/or a small dynamical range and/or inhomogeneous samples. Recent support for G00 results came from a study of $\\sim$ 200 galaxy groups, identified within the field galaxy redshift survey CNOC2 by Carlberg et al. (2001a), showing evidence that $M/L$ increases with increasing $\\sigma_v$. Moreover, $M/L$ of CNOC2 groups proves to be smaller than that of CNOC clusters (Hoekstra, Yee, \\& Gladders 2001). However, the question is still open. E.g., Hradecky et al. (2000), who computed homogeneous X--ray mass and optical luminosity for eight galaxy systems, claimed that $M/L_V$ is roughly independent of system mass. New insights on the behavior of $M/L$ for galaxy systems of different mass would be particularly useful in view of the theoretical predictions recently coming from cosmological N--body simulations combined with semianalytic modeling of galaxy formation. To draw more definitive conclusions about this topic, we extend the work of G00 by increasing the statistics of the data base and doubling the dynamical range, from $\\sim 5\\times 10^{13}$--$10^{15}$ \\msun to $\\sim 10^{12}$--$10^{15}$ \\msun. To this purpose, we consider both clusters analyzed by G98, the poor clusters by Ledlow et al. (1996; hereafter L96), the rich groups by Zabludoff \\& Mulchaey (1998a; hereafter ZM98), and the groups identified in the NOG sample (Nearby Optical Galaxy, Giuricin et al. 2000). The paper is organized as follows. We describe the data samples in \\S~2. We compute the main observational quantities, i.e. virial masses and optical luminosities, for all galaxy systems in \\S~3. We devote \\S~4 to the analysis of the relation between mass and luminosity and to the mass--to--light ratio. We discuss our results in \\S~5, while in \\S~6 we give a brief summary of our main results and draw our conclusions. Unless otherwise stated, we give errors at the 68\\% confidence level (hereafter c.l.) A Hubble constant of 100 $h$ \\ks $Mpc^{-1}$ is used throughout. ", + "conclusions": "" + }, + "0112/astro-ph0112372_arXiv.txt": { + "abstract": "The RICE experiment (Radio Ice Cherenkov Experiment) at the South Pole, co-deployed with the AMANDA experiment, seeks to detect ultra-high energy (UHE) electron neutrinos interacting in cold polar ice. Such interactions produce electromagnetic showers, which emit radio-frequency Cherenkov radiation. We describe the experimental apparatus and the procedures used to measure the neutrino flux. ", + "introduction": "Detection of ultra-high energy ($E_\\nu > 10^{15}$eV) neutrinos represents a unique opportunity to probe the distant universe. High-energy protons and photons from distant sources are likely to interact with the cosmic microwave background (CMB); protons, being charged, have their trajectories bent in galactic and intergalactic magnetic fields. Neutrinos are inert to CMB photons and point directly back to their source, giving essential information on those sources. In the realm of particle physics, detection of UHE neutrinos from cosmological distances, if accompanied by flavor identification, may permit measurement of neutrino oscillation parameters over an unprecedented range of $\\Delta m^2$. It has been suggested that, with a sensitive enough array, tau neutrinos may be identified by two-step ``double-bang\"\\cite{doublebang} processes where a $\\tau$ lepton is created and subsequently decays. Since neutrino absorption in the earth depends on the chord length through the earth,\\footnote{The Earth is opaque to $>$PeV neutrinos with Standard Model cross-sections and with zenith angles approaching 180 degrees. Largely because of earth shadowing, the RICE array is most sensitive to neutrinos incident at zenith angles between 60 and 120 degrees. Conversely, the angular distribution is sensitive to the cross-section and may allow checks of the Standard Model.} the angular distribution of detected neutrino events could be used to verify predictions for weak cross-sections at energies unattainable by any man-made accelerator. Alternately, if the high energy weak cross-sections are known, they can be used to test Earth composition models along an arbitrary chord (`neutrino tomography')\\cite{tomography}. \\subsection{Current Experimental Efforts} Several recent projects (including AMANDA\\cite{AMANDA99}, NESTOR\\cite{NESTOR}, Lake Baikal\\cite{BAIKAL99}, ANTARES\\cite{Antares00}) are optimized for detection of very high energy ($10^{12-15}$ eV) cosmic ray muon neutrinos. Sensitivities to higher energies, as well as electromagnetic cascades, have also been shown to be substantial in such experiments\\cite{TWR,ICRC-papers}. These instruments are based on photomultiplier tube detection of the optical Cherenkov cone from muons produced in muon neutrino charged current interactions. At high energies, muons have ranges of order 1 km and follow approximately straight trajectories (smeared by multiple scattering), punctuated by catastrophic bremstrahlung every 0.1--1 km or so, in which $\\sim$10\\% of the muon's energy is lost to a photon. RICE employs radio detection, which is believed to be the most efficient detection mechanism at energies of $10^{15}$ eV and beyond\\cite{Buford}. \\begin{figure}[htpb] \\begin{picture}(200,250) \\special{psfile=RICEcone.ps voffset=0 hoffset=120 vscale=50 hscale=50 angle=0} \\end{picture} \\caption{Simulated RICE event. The actual detector geometry is shown, to scale.}\\label{fig:evdisp} \\end{figure} The RICE concept is illustrated in Figure \\ref{fig:evdisp}, depicting a Cherenkov cone fit to a set of ``struck'' dipole receivers in a simulation event, along with the extracted neutrino direction. Receiver locations are drawn to scale in the Figure. In the actual array geometry, dipole receivers are spread over a 200 m $\\times$ 200 m $\\times$ 200 m cube beneath and around the Martin A. Pomerantz Observatory (MAPO), approximately 1 km. from the geographic South Pole. \\subsection{Radio Detection} We have initiated a pilot program based on radiowave receiver technology in order to extend electron neutrino detection up to the PeV energy scale. The detector is intended to have sensitivity in the energy range $E_{\\nu_e}\\sim 10^{15}-10^{18}$ eV. RICE (``Radio Ice Cherenkov Experiment'') should ultimately be capable of observing the sky with angular resolution of $\\sim$10-50 mrad. Coherent radio Cherenkov emission is an efficient method for detecting high energy particles. The history of the effect goes back to Jelley, who first considered whether cosmic ray air showers might produce a radio signal\\cite{Allan}. Askaryan\\cite{Askaryan} subsequently predicted a net charge imbalance in air showers, and coherent radio power proportional to the energy of the shower squared. Substantial radio emission from atmospheric electromagnetic cascades was observed more than 30 years ago\\cite{Allan,Jelley}. Progress in ultra-high energy air showers has sparked renewed interest, and new observations of radio pulses have been reported recently\\cite{Rosner}, suggesting a possible radio component to the Auger detector\\cite{Auger00}. A recent international meeting highlights current progress\\cite{RADHEP2000}. This effect has recently been observed in a test-beam experiment at SLAC\\cite{testbeam}. A beam of electrons, of known amperage, was fired into a sand target and the radiation resulting from the impact measured in the radio regime. As Figure \\ref{slac_fig4.ps} illustrates, the measured signal strength (diamonds) displays the expected dependence on both total beam current (left, solid curve) and frequency (right, dashed curve). \\begin{figure}[htpb] \\centerline{\\includegraphics[width=10cm]{slac_fig4.ps}} \\caption{Results of testbeam experiment[14], showing signal strength and expected dependence on total beam current (circles, left) and frequency (diamonds, right). Figure reprinted courtesy of D. Saltzberg and P. Gorham.} \\label{slac_fig4.ps} \\end{figure} \\subsection{Initiation of the RICE Experiment} The RICE experiment was initiated in October, 1995 when the AMANDA collaboration graciously consented to co-deployment of two shallow radio receivers (``Rx'') in the first holes being drilled for AMANDA-B. Following deployment, a surface transmitter (``Tx'') was used to verify that signals could be detected under-ice with better than 10 ns timing precision. However, cross-talk and amplifier oscillation problems precluded use of those receivers for science. The first three dedicated RICE receivers were deployed in 1996-97, along with one underice transmitter. The extraordinary RF clarity of South Polar ice was amply illustrated by the evident brightness of the AMANDA photomultiplier tubes 2 km. below the RICE receivers (and close to the null in the dipole antenna's reception pattern). A Fourier analysis of the RF transients produced by AMANDA photo-tubes in laboratory conditions indicated that PMT background power dominated at frequencies below 100 MHz, motivating use of high-pass filters in subsequent Rx deployments. Antenna deployments followed in 1997-98 (three more receivers and two more transmitters deployed in AMANDA-B holes), 1998-99 (six receivers deployed in 5''-diameter dedicated RICE holes, bored using a mechanical, rather than a hot-water drill), and 1999-2000 (six receivers, and one transmitter deployed in AMANDA-B holes). ", + "conclusions": "Basic calibration of both the time and amplitude response of RICE radio receivers has been made, relying primarily on data taken {\\it in situ}. A Monte Carlo simulation has been written which reproduces the gross features of those calibration data. The calibration of the detector is sufficient to allow limits to be placed on the incident high-energy neutrino flux. \\vspace{2cm}" + }, + "0112/hep-ph0112009_arXiv.txt": { + "abstract": "Electroweak baryogenesis could be very efficient at the end of an electroweak-scale inflation. Reheating that followed inflation could create a highly non-equilibrium plasma, in which the baryon number violating transitions were rapid. In addition, the time-dependent motions of the scalar degrees of freedom could provide the requisite CP violation. If the final reheat temperature was below 100~GeV, there was no wash-out of the baryon asymmetry after thermalization. The observed value of the baryon asymmetry can be attained in a number of models, some of which do not require a significant departure from the Standard Model. ", + "introduction": "Inflation probably took place in the early universe. If the scale of (the latest) inflation was of the order of the electroweak scale, the subsequent reheating would provide ideal conditions for electroweak baryogenesis. The original scenario of Kuzmin, Rubakov, and Shaposhnikov~\\cite{krs,rs} was based on a brilliant idea that all three Sakharov's conditions~\\cite{sakharov} necessary for a successful baryogenesis were satisfied, at least qualitatively, in the early universe due to the properties of the Standard Model at finite temperature. First, the baryon number was violated by sphaleron transitions. Second, the universe was out of thermal equilibrium during a phase transition. Finally, CP was broken, because it is not an exact symmetry in the Standard Model. Unfortunately, this scenario, in its original form, cannot account for the observed baryon-to-photon ratio $\\eta \\sim 10^{-10}$. First problem has to do with the fact that the electroweak phase transition is not sufficiently strongly first-order unless the Higgs mass is below about 45~GeV, which is ruled out by experiment. Second, the CP violation from the Cabibbo-Kobayashi-Maskawa matrix is way too small for baryogenesis because its contribution to $\\eta$ is suppressed by high powers of the Yukawa couplings. At the end of inflation, however, the universe was emphatically out of equilibrium. Could the baryon number violating processes take place in that non-thermal environment? The answer is yes; this was clearly demonstrated by a consensus of numerical~\\cite{ggks} and analytical~\\cite{kt,ggks,ck} arguments. Finally, the requisite CP violation may come from time-dependent solutions for the scalar zero-modes in the background~\\cite{ck}. Altogether, electroweak baryogenesis can be very efficient at the end of inflation. ", + "conclusions": "A highly non-thermal state of the universe at the end of inflation provides a fertile ground for baryogenesis. In addition to a dramatic departure from thermal equilibrium, preheating creates conditions for baryon number non-conservation and activates sources of CP violation that are largely unconstrained by the experimental data. Several modest modifications of the Standard Model allow for efficient electroweak baryogenesis in the wake of inflation." + }, + "0112/astro-ph0112366_arXiv.txt": { + "abstract": "Galactic HI is a gas that is coupled to magnetic field because of its fractional ionization. Many properties of HI are affected by turbulence. Recently, there has been a significant breakthrough on the theory of magnetohydrodynamic (MHD) turbulence. For the first time in the history of the subject we have a scaling model that is supported by numerical simulations. We review recent progress in studies of both incompressible and compressible turbulence. We also discuss the new regime of MHD turbulence that happens below the scale at which conventional turbulent motions get damped by viscosity. The viscosity in the case of HI is being produced by neutrals and truncates the turbulent cascade at parsec scales. We show that below this scale magnetic fluctuations with a shallow spectrum persist and point out to a possibility of the resumption of the MHD cascade after ions and neutrals decouple. We discuss the implications of the new insight into MHD turbulence for cosmic ray transport and grain dynamics. ", + "introduction": "The interstellar medium (ISM) is clumpy and turbulent (Larson 1981; Myers 1983; Scalo 1987; see also Lazarian, Pogosyan \\& Esquivel, this volume, henceforth LPE02) with an embedded magnetic field that influences almost all of its properties. This turbulence that ranges from AUs to kpc (Armstrong et al.~1995, Stanimirovic \\& Lazarian 2001, Deshpande et al.~2000) holds the key to many astrophysical processes (e.g., star formation, fragmentation of molecular clouds, heat and cosmic ray transport, magnetic reconnection). All turbulent systems have one thing in common: they have large ``Reynolds number\" ($Re\\equiv LV/\\nu$; L=characteristic size of the system, V=velocity different over this size, and $\\nu$=viscosity), the ratio of the time required for viscous forces to slow it appreciably ($L^2/\\nu$) to the eddy turnover time of a parcel of gas ($L/V$). A similar parameter, the ``magnetic Reynolds number\", $Rm$ ($\\equiv LV/\\eta$; $\\eta$=magnetic diffusion), is the ratio of the magnetic field decay time ($L^2/\\eta$) to the eddy turnover time ($L/V$). The properties of the flows on all scales depend on $Re$ and $Rm$. Flows with $Re<100$ are laminar; chaotic structures develop gradually as $Re$ increases, and those with $Re\\sim10^3$ are appreciably less chaotic than those with $Re\\sim10^7$. Observed features such as star forming clouds are very chaotic with $Re>10^8$ and $Rm>10^{16}$. {}Let us start by considering incompressible hydrodynamic turbulence, which can be described by the Kolmogorov theory (Kolmogorov 1941). Suppose that we excite fluid motions at a scale $L$. We call this scale the {\\it energy injection scale} or the {\\it largest energy containing eddy scale}. For instance, an obstacle in a flow excites motions on the scale of the order of its size. Then the energy injected at the scale $L$ cascades to progressively smaller and smaller scales at a rate of eddies turning over, i.e. $\\tau_l^{-1}\\approx v_l/l$, with the energy losses along the cascade being negligible\\footnote{This is easy to see as the motions at the scales of large eddies have $Re\\gg 1$ and therefore the energy loss into heat is negligible over the eddy turnover time.}. Ultimately, the energy reaches the molecular dissipation scale $l_d$, i.e. the scale where the local $Re\\sim 1$, and is dissipated there. The scales between $L$ and $l_d$ are called the {\\it inertial range} and it typically covers many decades. The motions over the inertial range are {\\it self-similar} and this provides tremendous advantages for theoretical description. The beauty of the Kolmogorov theory that it does provide a simple scaling for hydrodynamic motions. If the velocity at a scale $l$ from the inertial range is $v_l$, the Kolmogorov theory states that the kinetic energy ($\\rho v_l^2\\sim v_l^2$ as the density is constant) is transferred to next scale within one eddy turnover time ($l/v_l$). Thus within the Kolmogorov theory the energy transfer rate ($v_l^2/(l/v_l)$) is scale-independent, and we get the famous Kolmogorov scaling~~~~~$v_l \\propto l^{1/3}$~~~. One-dimensional\\footnote{Dealing with observational data, e.g. in LPE02, we deal with three dimensional energy spectrum $P(k)$, which, in isotropic turbulence, is related to $E(k)$ in the following way: $E(k)=4\\pi k^2 P(k)$.} energy spectrum $E(k)$ is one of the most important quantities in turbulence theories. Note that $E(k) dk$ is the amount of energy between the wavenumber $k$ and $k + dk$. When $E(k)$ follows a power law, $kE(k)$ is the energy {\\it near} the wavenumber $k\\propto 1/l$. Since $v_l^2$ represents a similar energy, $v_l^2 \\approx kE(k)$. Therefore, Kolmogorov scaling entails ~~~~~~$E(k) \\propto k^{-5/3}$~~~. Kolmogorov scalings constitute probably the major advance of the microscopic turbulence theory of incompressible fluids. They allowed numerous applications in different branches of science (see Monin \\& Yaglom 1975). However, astrophysical fluids are magnetized and the application of the Kolmogorov scalings is not easy to justify. For instance, dynamically important magnetic field should interfere with eddies motions. Paradoxically, astrophysical measurements reveal the Kolmogorov spectra (see LPE02). For instance, interstellar scintillation observations indicate electron density spectrum follows a power law over 7 decades of length scales (see Armstrong et al.~1995). The slope of the spectrum is very close to $-5/3$ for $10^6 m$ - $10^{14} m$. At larger scales LPE02 summarizes the evidence of $-5/3$ velocity power spectrum over pc-scales in HI. Solar-wind observations provide {\\it in-situ} measurements of the power spectrum of magnetic fluctuations and Leamon et al. (1998) also obtained a slope of $\\approx -5/3$. Is this a coincidence? What properties the magnetized compressible ISM is expected to have? These sort of questions we will deal with below. Here we describe our approach which is complementary to that in Vazquez-Semadeni (this volume). The latter attempts to simulate the ISM in its complexity by including many physical processes (e.g. compression, self-gravity) simultaneously. As a downside of this, such simulations cannot distinguish between the consequences of different processes. We discuss a focused approach when only after obtaining clear understanding on the simplest level, the next level is attempted. Therefore, we first consider incompressible MHD turbulence (\\S2), then discuss the viscous damping of incompressible turbulence in \\S3, and then we consider the effect of compression in \\S4. We discuss implications of our new understanding of MHD turbulence for the problems of dust motion and cosmic ray scattering in \\S5. ", + "conclusions": "Recently there have been significant advances in the field of MHD turbulence: 1. The first self-consistent model (GS95) of incompressible MHD turbulence that is supported by both numerical simulations and observations has been suggested. The major predictions of the model are scale-dependent anisotropy ($k_{\\|}\\propto k_{\\perp}^{2/3}$) and a Kolmogorov energy spectrum ($E(k)\\propto k^{-5/3}$). 2. Simulations of compressible MHD turbulence show that there is a weak coupling between Alfven waves and compressible MHD waves and that the Alfven modes follow the Goldreich-Sridhar scaling. {}Fast modes, however, decouple and exhibit isotropy. 3. On the contrary to the general belief, in typical interstellar conditions, magnetic fields can have rich structures below the scale at which motions are damped by viscosity created by neutrals (ambipolar diffusion damping scale). These advances change a lot in our understanding of many fundamental interstellar processes, e.g. cosmic-ray propagation and grain dynamics. In terms of HI they show a way to account for the formation of structures at very small scales. More discoveries are surely to come!" + }, + "0112/astro-ph0112199_arXiv.txt": { + "abstract": "We present an analysis of a large sample of moderate resolution Keck LRIS spectra of subgiant ($V \\sim 17.2$) and fainter stars in the Galactic globular cluster M5 (NGC 5904) with the goal of deriving C and N abundances. Star-to-star stochastic variations with significant range in both [C/Fe] and [N/Fe] are found at all luminosities extending to the bottom of the RGB at $M_V \\sim+3$. Similar variations in CH appear to be present in the main sequence turnoff spectra, but the signal in the current sample is too low for a detailed analysis. The variations seen among the M5 subgiants are consistent with the abundances found earlier by Briley \\etal\\ (1992) for brighter giants in this cluster. There is thus no sign of a change in the behavior of C and N with evolutionary stage over the full range in luminosity of the RGB and SGB, although a systematic decrease with luminosity in the mean [C/H] smaller than a factor of 2 cannot be ruled out with confidence at present. The C and N abundances appear strongly anti-correlated, as would be expected from the CN-cycle processing of stellar material. Yet the present stars are considerably fainter than the RGB bump, the point at which deep mixing is believed to set in. On this basis, while the observed abundance pattern is consistent with proton capture nucleosynthesis, we infer that the site of the reactions is likely not within the present sample, but rather in a population of more massive (2 -- 5$M_{\\odot}$) now defunct stars. The range of variation of the N abundances is very large and the sum of C+N increases as C decreases. To reproduce this requires the incorporation not only of CN but also of ON-processed material. Furthermore, the existence of this correlation is quite difficult to reproduce with an external mechanism such as ``pollution'' with material processed in a more massive AGB star, which mechanism is fundamentally stochastic in nature. We therefore suggest that although the internal mixing hypothesis has serious flaws, new theoretical insights are needed and it should not be ruled out yet. ", + "introduction": "} By virtue of their large populations of coeval stars, the Galactic globular clusters present us with a unique laboratory for the study of the evolution of low mass stars. The combination of their extreme ages, compositions and dynamics also allows us a glimpse at the early history of the Milky Way and the processes operating during its formation. These aspects become even more significant in the context of the star-to-star light element inhomogeneities found among red giants in every cluster studied to date. The large differences in the surface abundances of C, N, O, and often Na, Mg, and Al have defied a comprehensive explanation in the three decades since their discovery. Proposed origins of the inhomogeneities typically break down into two scenarios: 1) As C, N, O, Na, Mg, and Al are related to proton capture processes at CN and CNO-burning temperatures, material cycled through a region in the upper layers of the H-burning shell in evolving cluster giants may be brought to the surface with accompanying changes in composition. While standard models of low mass stars do not predict this ``deep mixing'', several theoretical mechanisms have been proposed (e.g., the meridional mixing of Sweigart \\& Mengel 1979, and turbulent diffusion, Charbonnel 1994, 1995) with varying degrees of success. Moreover, there is ample observational evidence that deep mixing {\\it{does}} take place during the red giant branch (RGB) ascent of metal-poor cluster stars (see the reviews of Kraft 1994 and Pinsonneault 1997 and references therein). 2) It has also become apparent that at least some component of these abundance variations must be in place before some cluster stars reach the giant branch. Spectroscopic observations of main sequence turn-off stars in 47 Tuc (Hesser 1978; Hesser \\& Bell 1980; Bell, Hesser, \\& Cannon 1983; Briley, Hesser, \\& Bell 1991; Briley \\etal 1994, 1996; Cannon \\etal\\ 1998) and NGC 6752 (Suntzeff \\& Smith 1991, Gratton \\etal\\ 2000), as well as our own work in M71 (Cohen 1999, Briley \\& Cohen 2001, Ram\\'{\\i}rez \\& Cohen 2002) have shown variations in CN and CH-band and Na line strengths consistent with patterns found among the evolved giants of these clusters. The assumption that these low mass cluster stars are incapable of both deep dredge-up and significant CNO nucleosynthesis while on the main sequence leads to the possibility that the early cluster material was at least partially inhomogeneous in these elements or that some form of modification of these elements has taken place within the cluster. Suggested culprits include mass-loss from intermediate mass asymptotic giant branch stars and supernovae ejecta (see Cannon \\etal\\ 1998 for an excellent discussion of these possibilities). Thus the observed light element inhomogeneities imply that there is some aspect of the structure of the evolving cluster giants which remains poorly understood (the deep mixing mechanism), that the early proto-clusters may have been far less homogeneous, that intermediate mass stars may have played a greater role in setting the composition of the present day low mass stars than previously thought, etc. Indeed, the evidence cited in the reviews above has led many investigators to suggest that a combination of processes is responsible, i.e., many clusters contain star-to-star inhomogeneities established early in their histories which have subsequently been further altered by deep mixing during the ascent of the RGB. This of course greatly exacerbates the difficulty of achieving an understanding of these issues, as a knowledge of the composition of the more easily observed bright red giants will not tell the whole story of their chemical history - one must also understand the makeup of the main sequence stars. In the present paper, we continue our earlier work on M71 by exploring the CH and CN band strengths in a sample of low luminosity stars in the somewhat more metal poor globular cluster M5. We adopt current values from the on-line database of Harris (1996) for the apparent distance modulus of M5 at $V$ of 14.31 mag with a reddening of E(B--V) = 0.03 mag. Recent CCD photometric studies of this cluster, focusing primarily on its age, are given by Johnson \\& Bolte (1998) and Stetson \\etal\\ (1999). Sandquist \\etal\\ (1996) discuss the predominantly blue horizontal branch of M5. We adopt the metallicity [Fe/H] = $-$1.21 dex found by Ivans \\etal\\ (2001) in a high dispersion abundance analysis of a large sample of stars on the upper giant branch of M5. We describe the sample in \\S\\ref{section_phot} and \\ref{section_spec}. We outline our measurement of the molecular band indices and their interpretation in terms of the scatter in \\S\\ref{section_indices}. With an assumption about the O abundance, these are converted into C and N abundances, from which we find a strong anti-correlation between C and N in \\S\\ref{section_cnabund}. A discussion of our results together with a comparison with the trends seen among the red giants in M5 and in other globular clusters is given in \\S\\ref{section_othergc} and \\ref{section_discussion}. A brief summary concludes the paper. ", + "conclusions": "} The primary facts that we have established for M5 are that there are strong stochastic variations from star-to-star of both C and N, with C and N anti-correlated, similar to those seen in M71. These variations are definitely present among the subgiants at the base of the RGB to $M_V {\\sim} +3$ and appear to extend to the main sequence stars as well. A question of considerable import for understanding the behavior of C and N in M5 is whether or not there is any change in the mean C abundance as one moves from the RGB tip down to the base of the subgiant branch such as is clearly seen in M15 and M92. As discussed in \\S\\ref{section_othergc} above, the best that can be said at the present time is that, unlike the case of M92 or M15, any systematic decline in the C abundance with evolutionary state as one moves up the RGB appears to be small, $<0.3$ dex. \\subsection{Implications for Stellar Evolution} A classical review of post-main sequence stellar evolution can be found in Iben \\& Renzini (1983). Their description of the consequences of the first dredge up phase, the only dredge up phase to occur prior to the He flash, indicates that a doubling of the surface N$^{14}$ and a 30\\% reduction in the surface C$^{12}$ can be expected, together with a drop in the ratio of C$^{12}$/C$^{13}$ from the solar value of 89 to $\\sim$20, as well as a drop in surface Li and B by several orders of magnitude. Observations of field stars over a wide range of luminosities conform fairly well to this picture, see e.g. Shetrone \\etal\\ (1993), Gratton \\etal\\ (2001), although additional mixing of Li and lower than predicted ratios of C$^{12}$/C$^{13}$ seem to occur even among field stars (do Nascimento \\etal\\ 2000). To match the observations of variations in abundances among globular cluster red giants which far exceed those described above, additional physics must be introduced into calculations of dredge up in old metal poor stars. Relevant phenomena include meridional mixing as described by Sweigart \\& Mengel (1979) as well as turbulent diffusion (see Charbonnel 1994, 1995) and the insights of Denissenkov \\& Denissenkova (1990) concerning the importance of the $^{22}$Ne($p,\\gamma)^{23}$Na reaction as a way to produce p-burning nuclei. The clear prediction of the most current calculations of this type by Denissenkov \\& Weiss (1996), Cavallo, Sweigart \\& Bell (1998) and Weiss, Denissenkov \\& Charbonnel (2000) is that the earliest that deep mixing can begin is at the location of the bump in the luminosity function of the RGB which occurs when the H-burning shell crosses a sharp molecular weight discontinuity. Zoccali \\etal\\ (1999) have shown that the luminosity of the RGB bump as a function of metallicity as determined from observation agrees well with that predicted by the theory of stellar evolution. Bono \\etal\\ (2001) further suggest that the agreement between the predicted luminosity function and actual star counts along the RGB in the vicinity of the bump in a suite of globular clusters is so good that mixing cannot have occurred any earlier, otherwise the evolutionary lifetimes, and hence the observed luminosity function, of such stars would have been affected by the mixing of He. Zoccali \\etal\\ (1999) give the expected location of the RGB bump in M5 to be 0.3 mag brighter than the HB, i.e. at $V \\sim14.8$ or $M_V \\sim+0.5$. Yet we see strong star-to-star variations in C and N abundances as well as a strong anti-correlation between them within a large group of cluster members at $V \\sim17.1~(M_V \\sim+2.8)$, more than 2.3 mags fainter at the base of the RGB. We see hints of such variation continuing on the upper main sequence at $M_V \\sim+3.7$. The range of luminosity over which these C and N variations are seen is becoming more and more of a problem for any scenario which invokes dredge up and mixing. Unless we have missed some important aspect of stellar evolution with impact on mixing and dredge up, we must declare the mixing scenario a failure for the specific case of M5 from our present work and M71 from our previous work (and several other globular clusters as well from the work of others). Even the theoreticians in the forefront of this field are beginning to admit that deep mixing alone is not sufficient (Denissenkov \\& Weiss 2001, Ventura \\etal\\ 2001). Unless and until some major new concept relevant to this issue appears, we must now regard the fundamental origin of the star-to-star variations we see in M5 as arising outside the stars whose spectra we have studied here. The strong anti-correlation between C and N, however, does suggest that CN-cycle material must be involved, and that this material has somehow reached the surface of these subgiant stars in M5. Since we know it cannot come from inside these stars, it must come from some external source. As reviewed by Lattanzio, Charbonnel \\& Forestini (1999), CN and ON cycling is known to occur in AGB stars, and AGB stars are also known to have sufficient dredge up to bring such material to their surfaces. We might speculate that the site of the proton exposure could be a previous generation of high mass stars, which then suffered extensive mass loss (either in or outside of binary systems) and polluted the generation of lower-mass stars we currently observe, while the higher mass stars are now defunct. \\subsection{ON Burning} Let us adopt as a working hypothesis that the C and N abundance variations we are seeing in the present subgiants are the result of the incorporation of material exposed to the CN-cycle (i.e., proton capture reactions) in a now evolved population of more massive (2 - 5$M_{\\odot}$) AGB stars, as was originally suggested in D'Antona, Chieffi \\& Gratton (1983). Indeed, recent models of metal-poor AGB stars by Ventura \\etal\\ (2001) suggest temperatures at the bases of the convective envelopes of such stars are capable of these reactions. The new generation of precision abundance analyses of globular cluster stars over a wide range of luminosity such as that of Ram\\'{\\i}rez \\& Cohen (2002) for M71 demonstrate that the abundances of the $\\alpha$-capture and s-process elements are constant, and so we further assume that the early cluster environment of M5 was not significantly polluted by the ejecta of even more massive stars. We now investigate whether this hypothesis is consistent with the C and N abundances we have derived for the main subgiant sample of M5. Figure~\\ref{fig_sum_cn} shows the sum of the C and N abundance as a function of the C abundance of the sample of M5 subgiants. The solid dot shows the predicted location assuming the initial C and N abundances (C$_0$, N$_0$) are the Solar values reduced by the metallicity of M5 ([Fe/H] = $-1.2$ dex). Thus this is the initial location for no burning and for a Solar C/N ratio. If the present stars incorporated material in which just C was burned into N, then the locus of the observed points representing the M5 subgiant sample should consist of a single horizontal line, with the initial point, the presence of no CN-cycle exposed material, at the right end of the line (the maximum C abundance) and the left end of the line corresponding to a substantial fraction of the star's mass (i.e. the atmosphere plus surface convection zone) including C-poor, N-rich AGB stellar ejecta. Furthermore, if the initial C/N ratio of the cluster is not Solar, then the locus should still be a horizontal line, but located at a different vertical height in this figure. The maximum possible N enhancement for a cluster SGB star with these assumptions occurs if the star formed entirely from AGB ejecta in which all C has been converted into N. For initial values (C$_0$, N$_0$) (not expressed as logarithms), this maximum N enhancement would be (C$_0$ + N$_0$)/N$_0$. If the initial value was the Solar ratio, C$_0$/N$_0 \\sim3.2$, the resulting maximum N enhancement is a factor of $\\sim$4.2, while for an initial C$_0$/N$_0$ of 10, the maximum N enhancement is a factor of 11. Now we examine the behavior of the C and N abundances among the M5 subgiant sample as inferred from our observations. It is clear that the assumption that the only thing happening is inclusion of material in which C was burned into N must be incorrect. The sum of C+N seems to systematically increase by a factor of $\\sim$5 between the most C rich star and most C deficient star. The discussion of the errors, both internal and systematic, in \\S\\ref{section_abund} suggest maximum systematic errors of $-0.2$ dex for log(C/H) and +0.2 for log(N/H). This is completely insufficient to explain such a large trend as errors. We thus have a serious discrepancy. The sum of C+N was {\\it{not}} constant as C was burned in the AGB sites. Furthermore the observed range in N abundances is very large. The most obvious way to reproduce this is to include O burning as well as C burning. If we adopt Solar ratios as our initial values, then a substantial amount of O burning is required. Figure~\\ref{fig_sum_cn} suggests that the initial ratio of C/N is not quite Solar, although not too far off. Adopting the Solar value as the initial C/N ratio, we calculate the minimum amount of O which must be burned at the base of the AGB envelopes to reproduce the locus observed in the figure (under the arguable assumption of the most extreme of our stars having formed largely from such material - this will, however, provide us with at least an estimate of the minimum burning required). We need to produce a N enhancement of a factor of 10. The Solar ratio is C/N/O = 3.2/1/7.6, so if all the C and 50\\% of the O were converted, we have an enhancement of N of a factor of 8 available to the present stars. Oxygen is typically found to be overabundant with respect to Fe in old metal-poor systems (see Mel\\'endez, Barbuy \\& Spite 2001, Gratton \\etal\\ 2001, Ram\\'{\\i}rez \\& Cohen 2002, and references therein); we assume [O/Fe] $\\sim +0.3$ dex, a typical value. Then the initial C/N/O ratios will be 3.2/1/15.2. Note that the same amount of O has to be burned to produce the observed distribution of C and N abundances, but in this case it is a considerably smaller fraction of the initial O. Returning to the AGB models of Ventura \\etal\\ (2001), the requirement for substantial O burning that emerges from our analysis of the CH and CN bands in the M5 subgiants may not be an unreasonable constraint - for metal-poor AGB stars they find temperatures sufficient for CNO-processing at the bases of AGB stars in a wide mass range. Following their Z=0.001 ([Fe/H] = --1.3) models we find surface O abundances dropping by a factor of 2 to 20 in masses from 4.5 to 2$M_{\\odot}$. We also note that under the assumption of little change in [C/Fe] (less than 0.3 dex as discussed above) taking place during the RGB ascent of the present low mass M5 stars, one should also expect little change in [O/Fe] as well, and that the O abundances of the present bright giants reflect their ``primordial'' values. The observed [O/Fe] abundances of the bright M5 giants by Sneden \\etal\\ (1992) are not inconsistent with this idea. Their ``O-rich'' stars average around [O/Fe] = +0.3 while their most ``O-poor'' stars are depleted by a factor of 3.5. If this is the result of ON-cycle exposure, more than enough N can be produced to explain the present results. This simple test is of course leaves significant questions unanswered. Problems include whether there were a sufficient number of AGB stars present to return the required quantity of material and the efficiency of any mechanism to incorporate it in the present stars. This is a non-trivial issue if the most C/O-poor, N-rich SGB stars formed with a preponderance of AGB ejecta - to reduce the C abundance by a factor of 6 by adding C-poor ejecta would require some 83\\% of the present star's mass to be made of this material. Note, as has been pointed out by several authors, these abundance inhomogeneities cannot simply be surface contaminations as they would be diluted by the increasing depth of the convective envelope during RGB ascent. Also, the range of C abundances among the M5 SGB stars appears much larger than that of any cluster studied to date. While this can perhaps be explained with regard to the more metal-rich clusters, whose polluting AGB stars should have undergone less ON-cycling, other even more metal-poor clusters than M5 appear to have smaller star-to-star C variations among their less evolved stars. This can be seen in the [C/Fe] versus luminosity diagrams of Carbon \\etal\\ (1982) and Bellman \\etal\\ (2001) for M92, where the range in [C/Fe] among the least luminous stars is relatively small (certainly not the factor of 6 seen here in M5). But this may be perhaps explained by varying the efficiency for incorporating AGB ejecta into subsequent generations of stars among the proto-globular clusters. Following Ventura \\etal\\ (2001), we predict substantial O variations, anti-correlated with Na and Al, to be present among the M5 subgiants and fainter stars. If the source of the proton exposed materials is indeed moderate mass AGB stars, these inhomogeneities should also correlate with Li abundances among the less evolved stars. We expect the verification or lack thereof of these predictions to be available shortly. Another important point is that Figure~\\ref{fig_sum_cn} shows no evidence for any bimodality in the distribution of the C and N abundances for the small section of the subgiant branch in M5 covered by our sample. The distribution along the locus of abundance appears to first order to be uniform. There is no preponderance for stars populating the extremes of high C or low C. An artificial suggestion of bimodality, or more correctly a tendency towards high CN strengths, could be produced by the distribution of C and N abundances shown in Figure~\\ref{fig_c_vs_n} or by saturation effects in the molecular bands themselves, which would only become apparent in cooler stars with stronger molecular bands. We also recall the anomalous stars in our sample, most of which are believed to be members of M5. Two of these in particular have enormously strong CH bands; we can offer no explanation for these stars at present. \\subsection{Additional Implications for Stellar Evolution} Figure~\\ref{fig_sum_cn} was used above to demonstrate that ON burning is required by considering the required N enhancement factor. This figure also displays a correlation between the sum of C and N number densities with the C abundance, i.e. the C abundance is correlated with the C/N ratio among the M5 subgiant stars. Here we explore the consequences of the existence of this correlation for the origin of the C and N variations themselves. Any external mechanism for producing these variations will involve an efficiency factor for the incorporation of material. We expect this factor to depend on the mass of the star itself, how much additional mass is incorporated (${\\Delta}M$), and the initial C and N abundances in the star itself and within ${\\Delta}M$. Since these properties of ${\\Delta}M$ might be expected to fluctuate wildly, this process therefore should show a lot of stochastic random variability. It is easy to imagine that various parcels of ${\\Delta}M$ have a wide range in C/N ratios due to the varying amount of nuclear processing that each might have experienced, and thus the strong stochastic star-to-star variations in C and N as well as their overall anti-correlation that we observe among the M5 subgiants can be reproduced. However, the correlation above requires a correlation between the mass of the parcel accreted, ${\\Delta}M$, and the C/N ratio of the material within this mass, and that seems to be rather artificial given the random nature of the process. We thus conclude that is difficult to reproduce the correlation described above with such an external mechanism involving accretion of ``polluted'' material from AGB stars. Unless we have made gross errors in the C and N abundances in our M5 subgiant sample far beyond what we believe might have occured, the existence of this correlation suggests that we should not totally rule out internal mechanisms as yet." + }, + "0112/astro-ph0112220_arXiv.txt": { + "abstract": "We present the results of a combined analysis of cosmic microwave background (CMB) and X-ray galaxy clusters baryon fraction to deduce constraints over 6 inflationnary cosmological parameters. Such a combination is necessary for breaking degeneracies inherent to the CMB. ", + "introduction": "Since their first detection by COBE, the CMB temperature fluctuations have become an essential tool for constraining cosmological parameters. From the beginning of 2000, new experiments have released data set of good quality up to the third acoustic peak. Better constraints have been obtained on several cosmological parameters. Nevertheless, it has been shown that even with precise measurements of the power spectrum, it is nearly impossible to distinguish models with the same physical parameters on the last scattering surface. Basically, some degeneracies are inherent to the CMB. We consider in this work X--ray clusters as an independent way for constraining cosmological parameters. ", + "conclusions": "Given the constraints of CMB and X-ray clusters, one could combined the two by multiplying the corresponding likelihood. \\begin{figure} \\plottwo{douspis_fig1.ps}{douspis_fig2.ps} \\caption{Constraints given by CMB and CMB plus X--ray data \\label{fig:contcmb}} \\end{figure} The right figure of fig. 1 shows the constraints from the combined analysis in the same plane as the left figure. We can see that the degeneracies are broken. Confidence intervals are now determined for $H_o$ and $\\Lambda$. For both parameters, low values are preferred: $25 < H_o < 50 $ and $\\Lambda < 0.45$ at 99\\% CL. Due to inherent degeneracies in the CMB it is nearly impossible to specify some of the cosmological parameter; ``cross constraints'' are then necessary. This work is thus a preliminary view of the kind of constraints one would obtain in a near future, using last X--ray satellite data, and in a less near future the CMB satellite data (MAP and Planck)." + }, + "0112/astro-ph0112016_arXiv.txt": { + "abstract": "In this contribution we schematically summarize a number of ongoing, coordinated spectro-photometric projects devoted to the study of the stellar population properties in $\\omega$~Cen. All of these investigations are part of a global comprehensive project led by our team to address the complex formation history of stars in this cluster. ", + "introduction": "After the recent discovery of a peculiar, extremely metal rich RGB population, we started a long term project devoted to the study of the multiple stellar populations in $\\omega$~Cen. The global project consists of two main branches: {\\it (a)} complete multi-wavelength photometric surveys to investigate the general properties of the sub-populations, from the far ultraviolet to the near infrared, using both wide-field and high-resolution imaging; {\\it (b)} high and moderate resolution, high signal-to noise, extensive spectroscopic campaigns for giant and subgiant stars both in the optical and the infrared wavelength ranges. As a final result, we expect to completely characterize the abundance patterns, the dynamical properties and the relative ages of the sub-populations in $\\omega$ Cen. A short description of some ongoing sub-projects is provided, along with some first preliminary results. ", + "conclusions": "" + }, + "0112/astro-ph0112293_arXiv.txt": { + "abstract": "Multiwavelength observations of the hard X--ray selected sources discovered by {\\it BeppoSAX}, {\\it Chandra} and XMM--{\\it Newton} surveys have significantly improved our knowledge of the AGN population. The increasing number of X--ray obscured AGN so far discovered confirms the prediction of those AGN synthesis models for the X--ray background based on the Unified scheme. However, follow--up optical observations of hard X--ray selected sources indicate that their optical properties are quite varied and the simple relations between optical and X--ray absorption are by no means without exception. Moreover there is evidence of a substantial number of luminous X--ray sources hosted by apparently normal galaxies. In this paper the results obtained from multiwavelength observations of hard X--ray selected sources discovered by {\\it BeppoSAX} and XMM--{\\it Newton} are presented and briefly discussed. ", + "introduction": "A large fraction of the energy density contained in the cosmic X--ray background spectrum (XRB) is accounted for by the summed contribution of Active Galactic Nuclei (AGN) if most of their high-energy radiation integrated over the cosmic time is obscured by gas and dust. Several independent models based on AGN unification scheme (Setti \\& Woltjer 1989, Madau et al. 1994, Comastri et al. 1995, Gilli et al. 2001) and energetic arguments (Fabian \\& Iwasawa 1999) lead to the conclusion that a fraction as high as 80--90\\% of the luminosity produced by accretion-powered sources is obscured at almost all wavelengths emerging only in the hard X--ray band above a few keV. Hard X--ray surveys represent thus the most efficient method to search for and to trace the cosmological evolution of accretion-powered sources. A still unknown fraction of obscured radiation is reprocessed and remitted in the far--infrared band. Multiwavelength follow--up observations of hard X--ray selected sources would allow to probe where the bulk of obscured accretion power is remitted and estimate the AGN contribution to the extragalactic background light in the infrared band. Thanks to their revolutionary capabilities (arcsec imaging and high-energy throughput) {\\it Chandra} and XMM--{\\it Newton} have opened up a new era in the study of the hard X--ray sky. Deep {\\it Chandra} surveys (Brandt et al. 2001, Rosati et al. 2001) have reached extremely faint fluxes in the 0.5--2 keV and 2--7 keV bands virtually resolving the entire XRB flux at these energies; while relatively deep XMM--{\\it Newton} exposures (Hasinger et al. 2001, Baldi et al. 2001) have extended by a factor 50 the sensitivity in the 2--10 and 5--10 keV bands with respect to previous {\\it ASCA} and {\\it BeppoSAX} observations. The X--ray source counts and their average spectral properties, which are now probed over several decades of fluxes and energy ranges, appear to be consistent with AGN synthesis model predictions. Although the remarkable results obtained so far, deep {\\it Chandra} and XMM--{\\it Newton} surveys are limited by small area coverage (less than a quarter of square degree) and by the extremely faint magnitudes of the optical counterparts which make the identification of X--ray sources very difficult, if not impossible, even at 8m class telescopes (Giacconi et al. 2001, Tozzi et al. 2001, Norman et al. 2001). In order to fully characterize the nature and evolution of the X--ray source population it is customary to complement deep surveys with shallower observations carried out on larger areas (see for example the Einstein Medium Sensitivity Survey: EMSS; Gioia et al. 1990). This approach allows us to minimize the effects of field-to-field fluctuations (the cosmic variance) and makes much easier the optical identification follow-up observations given the average brighter magnitude of the counterparts. In this review the results obtained by two Large area surveys carried out with {\\it BeppoSAX} ({\\tt HELLAS}) and XMM--{\\it Newton} ({\\tt HELLAS2XMM}) are summarized. The main scientific drive of this project is to probe the obscured accretion history of the X--ray Universe. The adopted strategy is a trade--off between the hardest X--ray band and the largest area which can be covered with {\\it BeppoSAX} and XMM--{\\it Newton}. ", + "conclusions": "The relatively high number of obscured AGN discovered by {\\it BeppoSAX} and XMM--{\\it Newton} makes high-energy large area surveys extremely well suited to study the physics and the evolution of the sources responsible for the hard X--ray background. The hard X--ray sky is populated by AGN with extremely varied broad-band properties. The most important result concerns the optical appearance of X--ray obscured AGN as a function of redshift. In the local Universe a fraction of absorbed objects are associated with apparently normal, early-type galaxies. The lack of any AGN feature in their optical--infrared spectra suggests the presence of buried, probably Compton-thick nuclei. At higher redshift the presence of hard X--ray sources with broad optical lines indicates that the absorbing gas is dust--free. Finally a sizeable fraction of hard X--ray selected sources lacks an optical counterpart at the limit of 4m class telescopes. Multiwavelength observations of hard X--ray selected sources allow us to uncover AGN activity in a number of objects which would have not been classified as such on the basis of observations at other wavelengths. Larger samples of hard X--ray selected sources will provide new insights into the physics and the cosmic history of accretion." + }, + "0112/astro-ph0112400_arXiv.txt": { + "abstract": "The HH\\,1/2 system has been observed with the two spectrometers on board the Infrared Space Observatory. Diffuse \\cii\\ emission indicate the presence of a Photo-Dissociation Region (PDR) which is only in part due to the external Far-UV irradiation from the nearby Orion Nebula. Additional irradiation must be originated locally and we show that the FUV field produced in the recombination regions behind the shocks traced by the HH objects is sufficient to induce a PDR at the flow cavity walls. The analysis of [O{\\sc i}], \\sii\\ and \\neii\\ lines suggests shock velocities v$_s\\sim$100$\\div$140 \\kms\\ with pre-shock densities between 100 and 1000 \\cmthree. H$_2$ pure rotational lines trace a T$\\sim$600 K gas which is likely to be warmed up in slow, planar, molecular shocks. The coexistence in the areas subtended by the instrumental beamsizes of shocks at different velocities is consistent with the bow-shock morphology of the HH\\,1 and, to a lesser extent, HH\\,2 as traced by optical images. ", + "introduction": "\\label{intro} The HH 1/2 is still the proto-typical optical Herbig-Haro (HH) outflow, and has been and it is the subject of large number of studies at almost all possible wavelengths, sensitivities and angular resolutions (see ``Notes'' in Reipurth~\\cite{bo99}~for nearly 100 references on this young stellar bipolar flow). Some of the most recent studies, however, have concentrated on the molecular gas properties and its relationship to the optical emission, particularly in the context of entrainment, shock chemistry, proper motions, multiple outflows, multiple ejection events and the 'unified model' for stellar jets-CO flows. It is in this realm that studies of the far-infrared emission are of great importance, since in these collisionally excited objects a large fraction of the atomic and molecular gas cooling takes place at these wavelengths (see e.~g. Saraceno~\\citeyear{sar98}; Molinari et al.~\\citeyear{Metal00}). A quick reminder of the HH 1/2 bipolar jet properties goes as follows; it lies in Orion at $\\sim 460$ pc, with an angular size of $\\sim 3\\arcmin$ on its brightest component but with signatures of shock excited emission from a previous ejection event $\\sim 23\\arcmin$ NW from VLA 1, the outflow's embedded Class 0 driving source (Ogura \\citeyear{ogu95}; Pravdo et al.~\\citeyear{prav85}). The proper motions indicate flow velocities of $\\sim 300-400$ \\kms\\ for both the atomic/ionic and warm H$_2$ gas (Herbig \\& Jones~\\citeyear{HJ81}; Eisl\\\"offel, Mundt \\& B\\\"ohm~\\citeyear{eis94}; Rodriguez et al.~\\citeyear{rod00}; Noriega-Crespo et al.~\\citeyear{nori97}), and this sets a dynamical age of $\\sim 6\\times 10^5$ yrs for the outflow, but with a major ejection event some $\\sim 7.5\\times 10^4$ yrs ago, defined by the brightest condensations. Recent single-dish radio observations have shown that the J=1-0 CO emission is bounded by the ionic/atomic gas (Moro-Martin et al.~\\citeyear{ama99}), thus providing strong support to the idea that molecular outflows are driven by the stellar jets themselves (see e.g. Raga~\\citeyear{raga94}). The Infrared Space Observatory (ISO, Kessler et al. \\citeyear{Ketal96}) observed HH 1/2 with the mid-IR camera (CAM) and the short and long spectrometers (SWS \\& LWS respectively). A partial analysis of circular variable filter (CVF) ISOCAM observations showed the presence of strong ground rotational H$_2$ emission lines with a morphology quite similar to that of the NIR H$_2$ gas. These data also showed some signatures of strong J-shocks, like \\neii, as expected in high excitation HH outflows (Cernicharo et al.~\\citeyear{cerni99}). In this study we present the analysis of the ISO SWS and LWS data with an emphasis on the properties of shocks in the HH 1/2 outflow and their effects on the immediately surrounding medium. A summary of the observations and data reduction techniques is presented in $\\S$\\ref{obs}, followed by a presentation of the results (\\S\\,\\ref{results}) and their discussion (\\S\\,\\ref{disc}). ", + "conclusions": "\\label{disc} \\subsection{The PDR} \\label{pdr} \\cii\\ line has similar fluxes, within 20\\%, along the flow; in particular, the line is stronger toward HH\\,2 and fainter toward VLA\\,1. On this basis, we exclude both the C-S star and the heavily embedded VLA\\,1 source as the FUV (between 6eV and 13.6eV) field sources responsible for the PDR conditions. Either the FUV field source is external to the HH\\,1/2 flow, or, if local, must be somehow distributed along the flow. We will see below that a combination of the two possibilities best explains the observations. \\subsubsection{External irradiation} \\label{pdr-external} The external irradiation hypothesis is supported by the evidence (Mundt \\& Witt \\citeyear{MW83}) of a significant UV irradiation on the HH\\,1/2 area from the Orion nebula, whose center lies $\\sim1.5$\\adeg\\ to the north. This contamination would amount to 30$\\div$50\\% of the flux levels detected toward HH\\,1/2 shortward of $\\sim$1500\\AA, decreasing to $\\sim$20\\% at longer wavelengths. Based on the IUE measured spectra (B\\\"ohm-Vitense et al. \\citeyear{Betal82}; B\\\"ohm, Noriega-Crespo \\& Solf~\\citeyear{BNS93}) the integrated FUV field due to irradiation from the Orion Nebula would amount to G$_0\\sim$7, although it should be noted that the data used by \\cite{MW83} do not go shortward of $\\sim 1500$\\AA. This value is compatible with the FUV field from the dominant Trapezium star ${\\rm \\theta}_c^1$ (about G$_0\\sim18$), with a reasonable amount of extinction along the path to HH\\,1/2. The maximum \\cii\\ line flux in the ISO-LWS beam (assuming complete beam filling) that can be produced by this FUV field is $\\sim3.5\\;10^{-19}$ \\lflux\\ for a gas density n$\\sim10^3$ \\cmthree\\ (Kaufman et al. \\citeyear{Ketal99}). The \\cii\\ flux decreases by about 40\\% for densities a factor 10 higher or smaller, and is almost linear with G$_0$. We conclude that about half, at most, of the PDR \\cii\\ emission measured by ISO-LWS may be due to external irradiation by the Orion Nebula. \\subsubsection{Local irradiation} \\label{pdr-local} In a recent paper (Molinari, Noriega-Crespo \\& Spinoglio \\citeyear{MNS01}) we proposed that the \\cii\\ flux detected toward the HH objects in the HH\\,80/81 flow was emitted in a PDR situated at the walls of the cavity excavated by the flow, and irradiated by collisionally-enhanced 2-photon FUV continuum emitted by the post-shock ionized regions in the HH objects (Dopita, Binette \\& Schwartz \\citeyear{DBS82}). We propose that this scenario is also pertinent to the HH\\,1/2 system. In terms of observables, this model establishes a relationship between the free-free radio continuum emitted by the recombination region in the HH object, and the \\cii\\ line flux radiated by the PDR at the cavity walls. In the optically thin regime, the former can be expressed as (Curiel, Cant\\'o and Rodr\\'{\\i}guez \\citeyear{CCR87}): \\begin{equation} S_{\\nu} = 1.84\\cdot 10^{-4} ~\\theta^2~ \\left[{{\\nu}\\over {10~ GHz}}\\right]^{-0.1}~ T_4^{0.45}~ n_{o_{10}} ~v_{s_7}~ [1+3.483 v_{s_7} -2.745]~{\\rm mJy} \\label{radioshock} \\end{equation} $\\theta$ is the angular diameter of the recombination region and is estimated off-the-plot from the 6 cm VLA D-conf. maps of \\cite{Retal90} to be $\\sim$10\\asec. T$_4$ is the electrons temperature in units of 10$^4$ K and we hold it fixed at 1 in these units. $v_{s_7}$ is the shock velocity in units of 100 \\kms\\ and we fix it to 1 in these units as determined in \\S\\,\\ref{atomic} below. $n_{o_{10}}$ is the pre-shock density in units of 10 \\cmthree\\ and we will keep this as the only free parameter in Eq.(\\ref{radioshock}). The \\cii\\ line flux can be expressed as (Molinari, Noriega-Crespo \\& Spinoglio \\citeyear{MNS01}): \\begin{equation} F_{C{\\sc ii}} = {{10^{-7}}\\over {D^2}} {{4\\pi}\\over{3}} r_{HH}^3 \\; \\eta_{ce} \\; f_c \\;\\chi_c \\int_{6eV}^{13.6eV} j_{\\nu} d\\nu~{\\rm W\\,cm^{-2}} \\label{fcii} \\end{equation} D is the distance of the source from the Sun. j$_{\\nu}$ is the 2-photon process emissivity and is a function, among other things, of the square of gas density. f$_c$ is the fraction of the 2-photon continuum emitted by the HH object which is intercepted by the portion of flow cavity wall encompassed by the instrumental beam used to measure the \\cii\\ flux, and can be derived with simple geometrical considerations. The radius of the HH object is set to 5\\asec\\ (see above). A critical parameter in Eq.(\\ref{fcii}) is the factor $\\eta_{ce}$ which accounts for the collisional enhancement of the hydrogenic 2s level population above the values predicted by pure recombination. $\\eta_{ce}$ can be determined from a comparison of the predicted two-photon spectrum with the observed UV continuum from HH objects. \\cite{DBS82} find values $5.5\\leq\\eta_{ce}\\leq10.6$ for the different knots of HH\\,1, and $2.8 \\leq \\eta_{ce} \\leq6.2$ for the different knots of HH\\,2. Since most of the knots in the two pointings will be encompassed by the LWS beam, we will adopt the mean $\\eta_{ce}$ values. Another important parameter is the fraction $\\chi_c$ of the incident FUV field which is re-radiated via the \\cii\\ line; this parameter is predicted to vary between 0.1\\% and 1\\% (Tielens \\& Hollenbach \\citeyear{TH85}). Since the rest of the FUV flux is reprocessed by dust into Far-IR continuum, our data should in principle allow us to independently estimate $\\chi_c$ as the ratio between the \\cii\\ line flux and the integrated FIR continuum as revealed by the LWS. Indeed, LWS observations clearly reveal significant continuum emission longward of $\\sim$50\\um\\ toward all pointings. The full spectra are shown in Fig. \\ref{lws_continua_wcm2um}, where the minispectra from the ten detectors have been stitched together for cosmetic purposes. The region is quite complex at these wavelengths and our continuum data do not represent an improvement, given the relatively poor spatial resolution, with respect to previous studies. Our data are in substantial agreement with \\cite{Cetal84}, also considering that our pointed positions do not match the emission peaks of the bipolar emission pattern mapped at 100\\um\\ with the KAO. The peak wavelengths in the SEDs of Fig. \\ref{lws_continua_wcm2um} imply black-body dust temperatures of 25, 32 and 47 K toward HH\\,2, VLA\\,1 and HH\\,2 respectively, in agreement with the KAO observations of \\cite{Hetal86}. The origin of the FIR continuum is not easy to assess. \\cite{Cetal84} do not detect 100\\um\\ emission from HH\\,2 to a 3$\\sigma$ limit of $\\sim$12Jy; clearly, a portion of the continuum that we detect toward HH\\,2 is due to contamination from VLA\\,1. As concerns HH\\,1 it is clear form the KAO data that a significant portion of the FIR continuum comes from a radio source $\\sim$30\\asec\\ S-SW of HH\\,1 and associated with an H$_2$O maser (Pravdo et al. \\citeyear{prav85}). Based on the same dataset, it seems also clear that the C-S star does not seem to significantly contribute longward of 50\\um; this is not surprising given its state of relatively evolved PMS star. \\placefigure{lws_continua_wcm2um} The pointing where a more reliable estimate of the PDR FIR continuum can be derived is probably HH\\,2. Indeed, the HH\\,1 continuum is severely contaminated by two sources (see above), while the VLA\\,1 continuum is likely dominated by thermal dust emission in the envelope of this Class\\,0 source. HH\\,2 instead, may only partially be contaminated by the VLA\\,1 continuum. The \\cii/FIR ratio toward HH\\,2 yields $\\chi_c\\geq0.005$, clearly in the range of values expected from the PDR models. With this choice of parameters we run our models (Molinari, Noriega-Crespo \\& Spinoglio \\citeyear{MNS01}) to determine the quantities in Eqs.(\\ref{radioshock}) and (\\ref{fcii}) as a function of the gas density n (in \\cmthree) and shock compression ratio $\\mathcal{R}$=n/n$_0$. The resulting grids are presented in Fig. \\ref{ciiradiofig} (full-line for HH\\,2, dashed-line for HH\\,1). The figure also shows the location of the HH objects based on the 6 cm flux observed by \\cite{Retal90}, and the \\cii\\ observed with the LWS, decreased by an amount $\\sim3.5\\;10^{-19}$ \\lflux\\ which is our estimate of the maximum contribution due to external FUV irradiation (\\S\\ref{pdr-external}). \\placefigure{ciiradiofig} The adopted model seems successful in reproducing the observable quantities. Formal values of density and compression ratio are [n, $\\mathcal{R}$]=[(5200,5100), (40,18)] for HH\\,1 and 2 respectively. A problem with this scenario, however, is that we seem to produce too much FUV continuum. The predicted \\emph{emitted} flux at the indicative wavelength of 1500\\AA\\ for the adopted model parameters is $\\sim 8\\,10^{-13}$ \\fdenscgs, which should be compared with the dereddened observed values of $\\sim10^{-13}$ \\fdenscgs\\ for the two HH objects (e.g., B\\\"ohm-Vitense et al. \\citeyear{Betal82}). We can tune the model parameters to produce less FUV continuum and still fit the \\cii\\ and 6 cm observations. For example, we may decrease the radii of the ionized regions by 50\\%, which is not unreasonable since we estimated this parameter directly from plots of radio maps. Additionally, we can also increase $\\chi_c$; in low irradiation conditions, like in the present case, this parameter can reach values $\\sim0.01$ (Tielens \\& Hollenbach \\citeyear{TH85}) which is not in contradiction with our previous estimate as this was a lower limit due to the difficulty of assessing the intrinsic FIR continuum of the HH objects (see above). With this choice of parameters we still fit the observed \\cii\\ line flux and 6 cm radio continuum, but with densities a factor 2 higher and compression ratios a factor 2 lower, producing an emitted 1500\\AA\\ continuum of $\\sim3\\,10^{-13}$ \\fdenscgs; the discrepancy with UV observations is reduced to a factor 3, which could be partly accounted for by the uncertainties in the reddenening corrections applied. As concerns the jet emanating from VLA\\,1, its low ionization (Solf \\& B\\\"ohm \\citeyear{SB91}) also confirmed by the low radio flux measured (e.g., Rodr{\\'{\\i}}guez et al. \\citeyear{Retal90}) is not sufficient to produce the amount of FUV continuum that would be required by our models to induce the \\cii\\ emission observed toward VLA\\,1. In the absence of significant dust column along the line of sight from the HH objects to the proximity of VLA\\,1, however, the direct FUV irradiation from the HH objects onto the cavity walls at the base of the flow would be sufficient to justify the \\cii\\ line flux observed toward VLA\\,1. The observation of a scattered-light component in the [S{\\sc ii}]$\\lambda\\lambda$6716,6731 line along the jet (Solf \\& B\\\"ohm \\citeyear{SB91}), and the presence of a biconical nebula seen in scattered optical light at the base of the jet (Strom et al. \\citeyear{Setal85}) would indeed seem to strengthen the idea of a relatively dust-free flow cavity (see also Davis, Eisl\\\"offel \\& Ray \\citeyear{DER94}). \\subsubsection{Line cooling} \\label{pdr-lines} We are now able to estimate the amount of cooling expected in lines other than \\cii\\ in the PDR conditions discussed above. The external and relatively faint FUV field (G$_0\\sim 7$) from the Orion Nebula can generate F(\\oia)$\\sim$5\\%F(\\cii), at most, for n$\\sim10^3$ \\cmthree. Increasing the density will also increase the expected \\oia/\\cii\\ cooling ratio (Kaufman et al. \\citeyear{Ketal99}), but the absolute \\cii\\ flux emitted would also decrease by a similar amount living the absolute \\oia\\ cooling essentially unaltered. As concerns the local FUV field, the G$_0$ value for the PDR will be dependent on the assumed distance between the HH objects and the flow cavity walls. The CO maps of \\cite{ama99} reveal a full width of the low-velocity lobes in the HH\\,1/2 outflow of the order of $\\sim30$\\asec, implying G$_0\\sim40$ at the cavity walls and a maximum \\oia\\ PDR cooling $\\sim20\\div30$\\% of the \\cii. As noted above, the \\emph{absolute} [O{\\sc i}] cooling is almost independent on the assumed density. Finally, high-J CO and rotational H$_2$ lines fluxes comparable to the observed values can only be generated in energetic (G$_0\\geq10^3$) and very dense (n$\\geq10^6$ \\cmthree) PDRs (e.g., Burton, Hollenbach \\& Tielens \\citeyear{BHT90}, which is clearly not the present case. \\subsection{Shock conditions in the HH objects} \\label{shock} \\subsubsection{Atomic lines} \\label{atomic} The shock conditions toward HH\\,1 and 2 were diagnosed comparing our observations (after correcting the [O{\\sc i}] lines for PDR contamination, see \\S\\ref{pdr-lines}) with the predictions from plane parallel atomic shock models. The shock models were calculated using MAPPINGS2, a code developed by \\cite{BDT85}, which has been thoroughly tested (Pequignot \\citeyear{P86}). A grid of models was created changing the pre-shock density between 10$^2$ and 10$^4$ \\cmthree, and the shock velocity between 60 and 140 \\kms. The rest of the model parameters have been fixed to standard values (Molinari, Noriega-Crespo \\& Spinoglio \\citeyear{MNS01}). The \\oia/\\oib$-$\\neii/\\sii\\ diagram in Fig. \\ref{hh1-2_diag} shows that shock velocities of v$_s\\sim$100 or $\\sim$140 \\kms, and pre-shock densities between 10$^2$ and 10$^3$\\cmthree\\ would be adequate to describe the observed line ratios. Higher shock velocities seem to be ruled out since \\neii/\\sii\\ ratios largely in excess with respect to the observations are produced (e.g., $\\sim$10 for v$_s\\sim$220 \\kms, using the planar shock models of \\cite{HRH87}). We note that the beamsize of the SWS at the \\sii\\ wavelength is a factor $\\sim1.7$ larger than at the \\neii\\ wavelength, but this would not change the position of the HH objects in Fig. \\ref{hh1-2_diag} by a significative amount even for uniformly extended emission. Noticeably, the estimated pre-shock densities are in good agreement with the independent estimates based on the shock-excited PDR model (\\S\\ref{pdr-local}). Those models were computed using v$_s$=100 \\kms, but we note that the adoption of v$_s$=140 \\kms\\ would simply imply, according to Eq. (\\ref{radioshock}), a right-shift of the model grids in Fig. \\ref{ciiradiofig}, still fitting the observations (with higher compression ratios, though). \\placefigure{hh1-2_diag} The comparison with optical line diagnostics is not straightforward given the complex morphology of the HH objects which appear structured in several knots, both in optical and near-IR images, which our observations do not spatially resolve (Herbig \\& Jones \\citeyear{HJ81}; Noriega-Crespo \\& Garnavich \\citeyear{NG94}; Davis, Eisl\\\"offel \\& Ray \\citeyear{DER94}). The SWS pointing toward HH\\,1 encompasses knots A, C, D and F; optical lines suggest shock velocities $\\sim80$\\kms\\ for A, and consistently above 130\\kms\\ for all the others (Hartigan, Raymond \\& Hartmann \\citeyear{HRH87}). The dust scattered emission requires at least 150 \\kms\\ shocks to explain the optical observations (Noriega-Crespo, Calvet \\& B\\\"ohm \\citeyear{NCB91}). Our estimate of v$_s\\sim100$\\kms\\ might represent average physical conditions over the SWS beamsize, although the negative detection of \\oiii\\ in the present observations (expected to be at the same flux level of the \\oia\\ line) cannot in principle be explained with v$_s\\geq$130\\kms. The presence of strong [O~III] 4959+5007 \\AA, however, in HH 1F and HH 2A+H (Solf \\& B\\\"ohm \\citeyear{SB91}) sets a lower limit at $\\sim 100$~\\kms\\ for the shock velocity and suggests that other mechanisms like, e.g., partial ionization of the pre-shock material, may be responsible for the depression of \\oiii\\ flux levels. Likewise, the SWS pointing toward HH\\,2 includes knots A, B, C, G, H and M, where shock velocities estimated from optical lines range from 100 to 200 \\kms~(Hartigan, Raymond \\& Hartmann \\citeyear{HRH87}; Solf \\& B\\\"ohm \\citeyear{SB91}). \\subsubsection{Molecular lines} \\label{molecular} Pure rotational H$_2$ lines in the ground vibrational level are normally analyzed under LTE conditions (e.g., Gredel~\\citeyear{G94}). Einstein coefficients and wavenumbers were taken from \\cite{BD76} and \\cite{D84}, and an ortho/para=3 has been assumed. The SWS aperture at the S(1) and S(2) line wavelengths is a factor $\\sim1.7$ larger compared with the other H$_2$ lines; fluxes for these two lines were corrected accordingly, assuming extended and uniform emission. We dereddened the H$_2$ line fluxes using the visual extinctions given by \\cite{Betal87} and using the \\cite{RL85} extinction curve. The H$_2$ gas temperature and column density can then be directly obtained from the slope and intercept of a linear fit to the Boltzmann plots shown in Fig.~\\ref{h2_excitation}, and are reported in Table \\ref{h2physparam} together with 1$\\sigma$ uncertainties in parenthesis; a solid angle equal to the SWS focal plane aperture for the S(3)-S(5) lines was assumed, i.e. 14\\asec\\ x 20\\asec~or 6.6$\\times10^{-9}$ sr (valid for all detected \\hii\\ lines but S(1), where the solid angle is $\\sim$30\\% higher), with a beam filling factor of 1. The fits show that a single temperature component is adequate to represent the observations, given the observational uncertainties. \\placefigure{h2_excitation} Using the same model and the parameters determined from the Boltzmann plots, we can estimate an H$_2$ cooling, for the component traced by the (0-0) lines, of $9\\;10^{-3}$ \\lsun\\ and $2.5\\;10^{-2}$ \\lsun\\ for HH\\,1 and HH\\,2, respectively. \\placetable{h2physparam} The issue of the nature of the H$_2$ excitation in HH\\,1 and 2 has been discussed by several authors in the last decade using near-IR imaging of the S(1)1-0 2.12\\um\\ and near-IR spectra in the H and K bands (Noriega-Crespo \\& Garnavich \\citeyear{NG94}; Davis, Eisl\\\"offel \\& Ray \\citeyear{DER94}; Gredel \\citeyear{G96}; Eisl\\\"offel, Smith \\& Davis \\citeyear{ESD00}; Davis, Smith \\& Eisl\\\"offel \\citeyear{DSE00}). The spectra suggest excitation temperatures ranging between 2000 and 3000 K in the various emission knots revealed in the images. Detailed modeling suggests that C-type (Draine \\citeyear{D80}) shocks best fit the observational scenario; H$_2$ emission would originate in the tails of the bow-shocks, whose tips are instead bright in [Fe{\\sc ii}]1.64\\um\\ (Davis, Eisl\\\"offel \\& Ray \\citeyear{DER94}), where the shock velocities are low enough for H$_2$ not to be dissociated. The fact that Boltzmann plots based on H$_2$ ro-vibrational are satisfactorily fitted with single linear fits led \\cite{G96} to doubt as to the C shock origin, since a stratification of temperatures up to $\\sim3000$ K should be expected. \\cite{ESD00} review the most popular shock models and indeed confirm that curved C-type shocks produce a quite shallow trend in the Boltzmann plots; besides, the predicted excitation temperatures of the $v=1$ lines are significantly lower ($\\sim1000$ K) than actually observed. Planar C-type shocks, on the other hand, seem to produce a rather abrupt change of slope at upper energy levels $\\sim4000-5000$ K. Since this is exactly at the separation between the $v=0$ lines we detected in the present work and the $v>0$ lines detected in the near-IR, the combined analysis of mid- and near-IR H$_2$ lines seems to suggest an origin in planar C-type shocks. C-type shocks are good emitters of molecular lines in general. Our marginal detection of high-J CO lines at flux levels comparable to H$_2$ (0-0) lines suggests shock velocities $15\\leq v_s \\leq 20$ \\kms\\ and pre-shock densities n$_0\\geq 10^5$ \\cmthree\\ (Kaufman \\& Neufeld \\citeyear{KN96}). However, for the same parameters the model also predict an H$_2$O(2$_{12}$-1$_{01}$) line flux about two orders of magnitude higher than actually observed in HH\\,2 (Tab. \\ref{restab}). Unless H$_2$O vapor is confined in a much smaller area compared to H$_2$ or CO, there is a clear deficit with respect to expectations. This seems unusual for Class-0 objects (Giannini, Nisini \\& Lorenzetti \\citeyear{GNL01}) and resembles more evolved Class-I systems. The possibility that water vapor depletes onto grains in the post-shock cooling region has been verified for HH\\,7 in NGC\\,1333 with the detection of the 62\\um\\ ice band (Molinari et al. \\citeyear{Metal99}); we do not detect such a feature anywhere in the HH\\,1/2 system, although the filling factor of the bow-shocks present in HH\\,1 and 2 (e.g., Fig. \\ref{hh1-2}) in the LWS beam is lower than in the case of HH\\,7 (also because of the larger distance). An interesting alternative to explain the water vapor deficit could be offered by the marginal detection of the fundamental lines of the OH molecule in HH\\,2. Given the high velocity of the dissociative shocks in HH\\,1/2, significant amounts of Far-UV flux can be expected to be irradiated by the post-shock regions (see \\S\\,\\ref{pdr}); in particular, the photons between 1300\\AA\\ and 2000\\AA\\ are capable of dissociating the water molecule into H and OH (Andresen, Thissen \\& Schroeder \\citeyear{ATS01}). Water photodissociation by a UV field was also invoked by \\cite{Setal00} to explain the surprisingly high OH cooling found toward T\\,Tau. Comparing the $v_s\\sim$20 \\kms\\ found for the C-type molecular shocks with the $v_s\\sim$100 \\kms\\ deduced for the dissociative J-type shock traced by the atomic lines (\\S\\,\\ref{atomic}), it is plausible that the two shock regions are respectively associated with the tails and tips of bow-shocks. Such structures are indeed identified in HH\\,1 and, with a more complex morphology, in HH\\,2 (Hester, Stapelfeldt \\& Scowen \\citeyear{HSS98}). The 2 orders of magnitude difference in pre-shock densities implies that significant compression of the pre-shock material has taken place in the molecular shocks. This is expected given the very nature of the C-type shocks, where the relatively low shock velocity and the presence of transverse magnetic field combine to accelerate and pre-compress the upstream material. I-band polarimetry (Strom et al. \\citeyear{Setal85}) indeed confirms the presence of a magnetic field oriented along the HH\\,1/2 flow." + }, + "0112/astro-ph0112383_arXiv.txt": { + "abstract": "We reconsider the rejuvenation mechanism as proposed by Shull, Fesen, \\& Saken (1989). These authors suggest that an active pulsar can catch up with, and rejuvenate the shell of the associated supernova remnant. The morphology of the SNRs G5.4-1.2 and CTB80 seem to confirm this rejuvenation mechanism. The spindown energy is deposited by the pulsar as a relativistic pulsar wind, and has a sufficient power to explain the observed radio emission observed in these remnants. Shull et al. (1989) did {\\it not} explain the observed lengthscales of the rejuvenated parts of the SNR shell. therefore one needs to consider the diffusive transport of the injected electrons by the pulsar wind. We propose to apply a diffusion mechanism as introduced by Jokipii (1987), which makes a distinction between diffusion along the magnetic field lines and perpendicular to the magnetic field lines, parameterised by the gyro factor $\\eta$. We show that one has to assume a high value for the gyro factor, $\\eta\\simeq 10^3-10^4$, i.e. diffusion of the electrons along the magnetic field line is much faster then perpendicular to the magnetic field line, in order for the rejuvenation mechanism to work on the observed lengthscales. ", + "introduction": "A supernova remnant (SNR) results from the supernova explosion of a massive star. When the core of the progenitor star collapses, a neutron star can be formed. Several mechanisms proposed for a core collapse supernova can impart a kick velocity to the pulsar, although no particular mechanism can be favoured. The relativistic pulsar wind will interact with the SNR, resulting in a pulsar wind nebula (PWN), which contains the shocked pulsar wind material. Initially the PWN will be centrally located in the SNR, but due to the kick velocity of the pulsar, the PWN will be dragged along by the pulsar, being deformed into a bow shock and ultimately break through the shell of the decelerating SNR (van der Swaluw, Achterberg, \\& Gallant 1998; Chevalier 1998). Figure 1 depicts the configuration of this last interacting stage of such a composite remnant. The relativistic pulsar wind is terminated by a strong MHD termination shock, whereas the bubble around the termination shock has been deformed into a bow shock due to the supersonic motion of the pulsar. Relativistic particles are injected at the site of the termination shock, of which a large fraction is advected away resulting in a wake of relativistic particles, which explains the observed trail of radio emission for bow shock PWN. However, due to the diffusive transport of the injected particles, a small fraction of these particles will radiate part of their energy away in the magnetic field of the SNR shell. Shull et al. (1989) have argued that these freshly injected particles brighten the radio emission of the SNR shell. We calculate the diffusive length scales of the injected electrons in the lifetime of a SNR or their radiative lifetime. This will enable us to compare these results with actual observed lengthscales in these systems. We will show that we need strongly anisotropic diffusion in order to obtain agreement between our calculations and the observed properties of several composite remnants. \\begin{figure*}[ht] \\begin{center} \\centerline{\\psfig{file=vdswaluwe_1.ps,height=3.5in}} \\end{center} \\caption{Configuration of a pulsar wind interacting with the shell of a supernova remnant. The pulsar wind is terminated by a strong MHD shock (dashed line) and the PWN itself is bounded by a bow shock (solid line). The PWN is propagating through the shell of its associated SNR, where the magnetic field lines are parallel with the shell of the SNR. By considering anisotropic diffusion, one cane explain the lengthscales of rejuvenated shells of SNRs like G5.4-1.2 and CTB80.} \\label{fig:1} \\end{figure*} ", + "conclusions": "We have reconsidered the rejuvenation mechanism as proposed by Shull et al. (1989). This was done by investigating the propagation of the injected relativistic electrons at the site of the wind termination shock through the PWN and the associated SNR shell. We conclude that a pulsar wind can rejuvenate the shell of a SNR if the diffusive transport of the injected electrons is strongly anisotropic, i.e. the gyro factor has to have a minimum value of $\\eta\\sim 10^3-10^4$. The limits for the gyro factor derived in this work are {\\it lower} limits because of the assumptions we made about the interaction time and the configuration of the magnetic field (section 2). Furthermore, our numerical experiment has shown that the line of sight between the observer and the magnetic field introduce a projection effect, which diminishes the lengthscales of the rejuvenated parts of the SNR shell (section 3)." + }, + "0112/astro-ph0112330_arXiv.txt": { + "abstract": "We use the radio ring in JVAS~B0218+357 to constrain the mass models with LensCLEAN. $H_0$ is determined from the resulting model and the time delay. B0218+357 is one of the best systems to reduce systematic errors in $H_0$ and may thus be called a ``golden lens''. ", + "introduction": "Gravitational lensing offers the unique possibility to determine the Hubble constant $H_0$ (Refsdal 1964) involving the understanding of very little astrophysics, keeping statistical and systematic errors well under control. The main uncertainty is introduced by the mass models of the lens. For multiply imaged compact sources, the number of constraints is very small. To overcome this difficulty, lensed extended sources can be used. ", + "conclusions": "" + }, + "0112/astro-ph0112106_arXiv.txt": { + "abstract": "\\noindent We present phase-coherent timing solutions obtained for the first time for 17 pulsars discovered at Arecibo by Hulse \\& Taylor (1975ab) in a 430-MHz survey of the Galactic plane. This survey remains the most sensitive of the Galactic plane at 430\\,MHz and has comparable equivalent sensitivity to the 1400-MHz Parkes multibeam survey. Comparing both surveys we find that, as expected, the one at 430\\,MHz is limited in depth by interstellar dispersion and scattering effects; and that the detection rate of pulsars with high spin-down luminosity ($\\dot E>10^{34}$\\,erg\\,s$^{-1}$) at the low frequency is a factor of 5 smaller than at high frequency. We also present scatter-broadening measurements for two pulsars and pulse nulling and mode-changing properties for two others. ", + "introduction": "\\label{sec:intro} The first systematic search for pulsars using the 305-m Arecibo radio telescope was carried out by Hulse \\& Taylor (hereafter HT) over 25 years ago \\nocite{ht74,ht75a,ht75b} (HT 1974, 1975a, b). The HT survey was conducted at 430\\,MHz and covered an area of about 140 square degrees in the region defined by $42\\arcdeg \\la l \\la 60\\arcdeg$ and $\\mid\\!\\!b\\!\\!\\mid\\:\\la 4\\arcdeg$. An integration time of 136.5\\,s resulted in a limiting flux density to long-period pulsars of about 1\\,mJy --- roughly an order of magnitude more sensitive than any other pulsar survey at that time, and still the most sensitive low-frequency survey along the Galactic plane. A total of 40 pulsars were discovered, the most notable being the double neutron star binary system B1913+16 \\nocite{ht75a,tw82,tw89} (HT 1975a) which has since been used as a magnificent laboratory for gravitational physics (Taylor \\& Weisberg 1982, 1989). Pulse period, position and dispersion measure (DM) determinations for the 39 long-period pulsars were published shortly after the completion of the survey (HT 1975b). However, perhaps as a result of the interest in observations of B1913+16, timing observations have been carried out for only 22 sources \\nocite{gr78} (see e.g.~Gullahorn \\& Rankin 1978). In this paper we report on new Arecibo timing observations of the remaining 17 pulsars which have resulted in accurate parameters for all of them. In \\S~2 we describe the observations and procedures used to obtain the timing solutions. The results are presented in \\S~3. We discuss scattering and single-pulse properties of the pulsars respectively in \\S~4.1 and \\S~4.2. Finally, in \\S~4.3, we briefly discuss the population properties of the sample and compare them with pulsars discovered in the 1374-MHz Parkes multibeam survey. ", + "conclusions": "" + }, + "0112/astro-ph0112276_arXiv.txt": { + "abstract": "We present a harmonic model for the data analysis of an all-sky cosmic microwave background survey, such as {\\it Planck}, where the survey is obtained through ring-scans of the sky. In this model, resampling and pixelisation of the data are avoided. The spherical transforms of the sky at each frequency, in total intensity and polarization, as well as the bright-point-source catalogue, are derived directly from the data reduced onto the rings. Formal errors and the most significant correlation coefficients for the spherical transforms of the frequency maps are preserved. A clean and transparent path from the original samplings in the time domain to the final scientific products is thus obtained. The data analysis is largely based on Fourier analysis of rings; the positional stability of the instrument's spin axis during these scans is a requirement for the data model and is investigated here for the {\\it Planck} satellite. Brighter point sources are recognised and extracted as part of the ring reductions and, on the basis of accumulated data, used to build the bright-point-source catalogue. The analysis of the rings is performed in an iterative loop, involving a range of geometric and detector response calibrations. The geometric calibrations are used to reconstruct the paths of the detectors over the sky during a scan and the phase offsets between scans of different detectors; the response calibrations eliminate short and long term variations in detector response. Point-source information may allow reconstruction of the beam profile. The reconstructed spherical transforms of the sky in each frequency channel form the input to the subsequent analysis stages. Although the methods in this paper were developed with the data processing for the Planck satellite in mind, there are many aspects which have wider implementation possibilities, including the construction of real-space pixelised maps. ", + "introduction": "The traditional method used to analyse the continuum distribution of radiation on the sky is through pixelisation: the calibrated measurements are projected onto pixels of a sky map, and further analysis of the data is done by means of the pixelised responses. These methods, which are most powerfully explored in the HEALPix software \\cite{G2}, have the advantage of providing a visual link to the reductions. Another already well-developed example of pixelisation is the IGLOO scheme developed by Crittenden \\& Turok~(1998) and Crittenden~(2000). A properly designed pixelisation scheme can, in addition, simplify the further analysis, allowing the use of fast spherical transform algorithms, and removing contaminated regions of the sky (G{\\'o}rski, 1994; Mortlock, Challinor \\& Hobson, 2001). A problem with any pixelisation of survey data is the loss of information due to the binning of the data. As a result of coverage variations, different pixels have different formal errors which complicates the application of fast spherical transforms. Pixels may be correlated with neighbouring pixels, which further complicates the analysis. The effective beam profile is unlikely to be circularly symmetric (as a result of intrinsic asymmetries as well as due to unconvolved sampling intervals) and requires the observations to be deconvolved. The effective two-dimensional beam profile applicable to a pixel therefore depends on the orientations of the scans which have contributed to it, information which is complicated to maintain. Thus, unless a range of `parallel maps' with supplementary information is kept (the implementation of which would do away with most of the advantages of the pixelisation), a pixelisation of the data will inevitably lead to loss of information. This loss affects our ability to recover the statistical properties of the reduced data, and as such should be considered a potentially serious disadvantage. The products of a survey mission like {\\it Planck} (see for example http://astro.estec.esa.nl/planck) or {\\it MAP} (http://map.gsfc.nasa.gov/m$\\_$mm.htm) are a point source catalogue and a set of frequency maps, from which maps of astrophysical components are derived. Methods for the component separation have been developed amongst others by Hobson et al.~(1998), Prunet et al.~(2001) and Baccigalupi et al.~(2000). In the latter paper an independent component analysis (ICA) algorithm is used to separate statistically independent signals from the frequency maps. In Hobson et al.~(1998) component separation on small rectangular fields is investigated using a maximum entropy technique in the Fourier domain. The latter method has been further developed into a full sky implementation at the resolution required for the {\\it Planck} survey \\cite{S2}, with the frequency and component maps analysed in the form of their spherical transforms. The analysis of the cosmic microwave background (CMB; one of the components separated) depends ultimately on the reliability of the separation process. A major aim of the data analysis should be to derive the input for the component separation, the spherical multipoles of the frequency maps, in the most reliable way, preserving as much information as possible on formal errors and their correlations. The analysis of data from an all-sky survey mission in which the survey is built up from circular scans can be carried out efficiently in three steps, as was done for the Hipparcos reductions (Lindegren, 1979; ESA, 1997): the first step reduces the raw measurements to ring data; the second step combines the ring data obtained over the mission to reconstruct the sphere; the third step analyses the sphere data. The {\\it Planck} mission fits this model very well: data is obtained over 1 hour intervals during which the spin axis of the satellite remains in a nominally `fixed' position. The collection of scientifically useful data starts at the end of a pointing manoeuvre and finishes with the start of the next manoeuvre, which moves the spin axis to the next scan position. During this interval, which is referred to as a time-ordered period (TOP), the satellite performs some 60 revolutions. The area of the sky covered by a detector during a TOP is referred to as a `ring'. The spinning of the satellite causes every detector $d$ to describe a small circle on the sky, each with its own specific opening angle $\\alpha_d$ (the opening angle is equivalent to the co-latitude of the ring as seen from the spin-axis position). A slight wobble of the spin axis (nutation) widens the rings, but this effect is small relative to the beam width. Data in a ring are referred to the circle defined by the mean spin-axis position of the satellite during the TOP and the opening angle of the detector concerned. The nutation effect will simulate a small-amplitude periodic variation of the opening angle for the detector, which will be most noticeable for point sources close to the rings as observed by the three highest frequency channels. Over a one-year mission almost 9000 rings will be produced for each of the 48 High Frequency Instrument (HFI) and 56 Low Frequency Instrument (LFI) detectors. In the harmonic model the three data analysis steps can be summarised as follows: \\begin{enumerate} \\item The analysis of the measured samplings per individual detector, collected over each TOP. The samples are referred to phases along reference circles of which the normal through the centre corresponds with the mean spin-axis position over the TOP. After cleaning these data from spikes (primarily resulting from very high energy radiation and particles), the phase-ordered data is collected in phase bins for further treatment. The short-term response variations are derived, using the accumulated differences between the mean responses in each phase bin and the individual contributions. After applying the response corrections, the mean sample values per bin are re-calculated and the data is ready for analysis. The first step in the analysis is the point-source transit identification. This is initially carried out on the phase-binned data, but in later iterations the point source catalogue constructed from the data is also used. The entire phase-binned signal is then fitted with Fourier components for the background signal and corrections to the assumed abscissae and intensities for all identified point sources. The corrected abscissae and intensities, with their formal errors, form the input to the point-source catalogue construction. The transit profiles of the brightest point sources are accumulated to provide the input for the central-beam-profile calibration. \\item Joint analysis of the Fourier components of the rings from all detectors operating at the same frequency, to produce the spherical transforms of the frequency maps. For each ring and each detector there is a unique coupling matrix, which describes how the spherical multipoles are projected on a small circle with a given opening angle $\\alpha_d$ and a two-dimensional beam profile. The accumulated coupling matrices and ring harmonics for a given frequency are the input to a generalised least squares solution for the multipoles of that frequency map. In an iterative process, the long-term response variations are calibrated; calibration of the outer beam profile may also be done at this stage. The processing up to this point is very similar for scalar and polarised data \\cite{C1}. \\item The analysis of the spherical multipoles of the frequency maps (for example, by means of the maximum entropy method) to extract the faint point sources, the Sunyaev-Zel'dovich clusters (Sunyaev \\& Zeldovich, 1980) and the spherical multipoles of the component maps, which are further analysed in the form of their power spectra. \\end{enumerate} Due to the way these steps are interlinked with geometric and response calibrations and the construction of the point source catalogue, they will require iterations before reduction results can be considered satisfactory. The present paper focuses on the the first of the three steps mentioned above. Although this is seen here as the first step in the harmonic model for the data analysis, many considerations made in the reductions as described here are equally applicable in a more traditional pixelised approach. The second step is described in the accompanying paper \\cite{C1}, and covers the reductions of the ring data to spherical multipoles of frequency maps in detail, a process which is more specific to the harmonic data model. The third step, described by Stolyarov et al.~(2002), is the full-sky maximum entropy component separation. Additional aspects concerning the analysis of incomplete sky maps are discussed in Mortlock et al.~(2002). An important criterion for the implementation of the harmonic model is the stability of the spin-axis pointing. This stability is affected by two main contributors: residual velocities around the two axes perpendicular to the spin axis, which will create nutation; and external torques, which can create residual velocities. The amplitudes of the residual velocities at the start of a TOP are determined by the criteria for the nutation damping, and should be small enough not to create a significant effect (in comparison with the beam width) on the accumulated data. The nutation wobble of the spin axis and the effect of the external torques can be derived from an analysis of the satellite's attitude in analytic form and through numerical integration. An analysis of this kind for the {\\it Planck} satellite is presented in Appendix~\\ref{App_att}, the results of which are summarised in Section~\\ref{Sec_att}. In Section~\\ref{Sec_calibr} we show how the data analysis in the harmonic model is closely linked with geometric and response calibrations. Details of the data analysis method are presented in Section~\\ref{Sec_circ}. Section~\\ref{Sec_point} describes the construction of the point-source catalogue and the calibration of the reference phases for the reduced TOPs. Section~\\ref{Sec_geom} presents the geometric calibrations: the focal plane geometry, the opening angle correction and the focal plane orientation correction. Finally Section~\\ref{Sec_iterate} describes the way the data reduction and calibration processes are linked in iterative loops. \\section[]{Attitude analysis for the \\textbfit{Planck\\/} mission} \\label{Sec_att} The results of a full analytic and numerical analysis of the {\\it Planck} satellite attitude are presented in Appendix~\\ref{App_att}. Here we summarize those aspects which are of immediate importance for the data analysis. Though this analysis is applied to parameters associated with the {\\it Planck} satellite, it can easily be modified to apply to different satellite configurations, as long as the outer product of the rotational velocity and angular momentum vectors is high with respect to the external torques acting on the satellite. The {\\it Planck} satellite is designed for a survey mission: over a period of a year the entire sky will be scanned twice, providing maps of the microwave sky at pass bands with central frequencies between 30~GHz and 857~GHz. The scanning will take place at pre-determined nominal pointings of the spin axis (pointing in a roughly anti-Solar direction). These nominal positions will not be exactly reproduced: a pointing noise of a few arcmin is expected. The satellite will rotate at a nominal velocity of 1~rpm, but small variations are expected here too. The interval of $\\sim 1$~hour between two re-positioning manoeuvres of the spin axis is referred to as a TOP. During a TOP, the satellite's spin axis will describe a relatively stable nutation ellipse in the satellite coordinate system (for as long as it remains unaffected by nutation damping), the maximum amplitude of which is determined by the residual velocities around the $x$- and $y$-axes, and should be less than 1.5~arcmin at the end of the nutation damping. The period of the nutation will be around 3~to 4~minutes, and will slowly change over the mission due to changes in the inertia tensor. The noise on the movement of the spin axis of the satellite relative to the nutation ellipse is at the arcsecond level. Such variations are not relevant for detectors with a 5~arcmin or more \\textsc{fwhm} response beam. This allows one to incorporate a simple set of nutation ellipse parameters in the data analysis if necessary. \\begin{figure} \\centerline{\\psfig{file=MB792_fig1.eps,width=8truecm}} \\caption{The orientation angles in the {\\it Planck} scan. The spin axis position (SA) is defined by its ecliptic coordinates $(\\lambda_z,\\beta_z)$, measured from the vernal equinox (VE); the spin axis position and the North Ecliptic Pole (NEP) define a meridian, called the `reference circle'; the position of a ring is defined by its opening angle $\\alpha_d$; the scan phase $\\psi$ defines the angle between the fiducial reference point (FRP) in the field of view (FOV) and the reference circle.} \\label{fig_ang} \\end{figure} The rotation of the satellite will cause every detector $d$ to describe a small circle on the sky, at an angular separation $\\alpha_d$ (referred to as the opening angle) from the position of the spin axis (see Fig.~\\ref{fig_ang}). As the actual spin axis will not coincide exactly with the nominal spin axis as defined in the hardware, there will be a difference between the actual opening angle and its hardware specification value. The wobble of the spin axis slightly widens the ring, but for most detectors the beam width is such that this can be ignored. This may not be the case for the highest frequency detectors, where the spin axis wobble may add a little to the beam width as observed perpendicular to the scan direction. The way it does so will be unique to each ring set due to the variable conditions at the end of nutation damping (see Appendix~\\ref{App_att}). We can conclude that, within the model presented, the attitude of the satellite can be described in a simple manner, while the actual attitude motions do not cause any problems for the harmonic analysis of the data. The chance of significant unaccounted forces acting on the satellite, which could disturb this situation, is small. \\section[]{The harmonic data model and system calibrations} \\label{Sec_calibr} The harmonic data model aims to provide an accurate analysis of the data and its statistical properties. It does so through a sequence of reduction steps, resulting in frequency maps in the form of spherical multipoles with their covariance matrices and a list of point source positions and intensities with associated formal errors. These maps can then be used as input for the component separation process, which can be performed in spherical multipole space \\cite{S2}. It should be possible to derive the spherical transforms of the frequency maps from the Fourier components of the rings, using information on pointing, focal plane geometry, focal plane orientation, the opening angle, the noise correlations of the data and the beam profile, although a number of computational problems makes this stage non-trivial \\cite{C1}. The {\\it Planck} mission is designed to obtain full-sky surveys at nine different wavelengths. It explores new ranges of resolution, coverage and sensitivity at these wavelengths. For such pioneering observations, calibrations are best done by demanding internal consistency of the observed sky, which in most cases will not have changed from one observation to the next. The exceptional variable and moving sources should all be identifiable. Demanding internal consistency defines simple, differential reduction procedures, and should result in mission products with the highest possible internal precision. Absolute calibrations are carried out only in the final stages of the data processing, using a selection of sources that can be assigned and measured in a well-controlled way with ground-based or other suitable experiments. It is important during the absolute calibrations to consider the spectral response of each source and its convolution with the {\\it Planck} passbands, and any background contributions which may be differently perceived by different instruments. The latter applies in particular to calibration sources close to the galactic plane. The preference for an internal, differential calibration applies to both responses and geometric properties. For the responses it involves the following steps: \\begin{enumerate} \\item The short-time response variations (time scales from one minute to one hour) are calibrated as part of the ring reductions, using the phase-binned data (see Section~\\ref{Sec_phase}). They show up by displaying, as a function of time, the differences between the mean samples per bin and the individual contributions to each bin. Removing any systematic effects observed in these residuals reduces all data contained in one TOP to an arbitrary mean response level for a specific detector and ring. All variations at shorter time-intervals are reduced to noise contributions, some of which may be correlated between neighbouring samples. This is taken into account in the noise correlation matrix for the model-fit solutions. \\item The long-term response variations (with time scales from one hour to one year) are calibrated as part of the sphere reductions. The reconstructed spherical transforms will, in the first instance, represent an average of all mean ring responses. By relating (in Fourier space) the response of each ring to this mean sky, response corrections per ring and detector are obtained. This is the equivalent of de-striping \\cite{R1}, though in the method proposed here all data on the ring will contribute. In case the instrument high-pass filters the signal, and there is essentially no information left on the monopole, these adjustments can still be done on the remaining signal. The response-corrected rings can be used again in a reconstruction of the spherical transforms, which should result in a decrease in the fitting noise. The process can be repeated until no further significant changes to the ring responses are observed. Given the expected large number of rings involved (9000 per detector), this process is expected to converge quite rapidly. The Hipparcos reductions have demonstrated the validity of this approach in the iteration between the great circle reduction, the sphere reduction and the improvements to the astrometric parameters, which was based entirely on internal consistency of the data (see for example van Leeuwen, 1997). External information may be used to obtain a reasonable first estimate of the calibration parameters to ensure a rapid convergence of the process. \\item The last step in the response calibration adjusts each final spherical transform of a frequency map using a set of ground-based (or other) calibration sources (see above), thus transforming them from relative to absolute scales. \\end{enumerate} The result of these processes is a set of frequency maps whose {\\it precision} is entirely determined by the internal calibrations, and {\\it accuracy} by the quality and quantity of external calibration sources. At any stage, should improvements in external sources become available, the external calibration can be performed again. In addition, as no external information is used in the internal calibrations, the loss of information from the original data is minimal. The loss is determined primarily by the limitations in the modelling of detector response variations. If instead, for example, properties of the dipole are used in the earlier calibration steps, the extraction of independent information on the dipole from the {\\it Planck} data could be compromised. The accuracy of the calibration parameters would depend on the local amplitude of the dipole signal as projected on a ring, which varies strongly with the ecliptic longitude of the spin axis. The response or beam profile forms the link between the detector response and the detector pointing. The calibration of the central part of the beam profile is closely linked with the geometric calibrations and the construction of the bright point source catalogue. The geometric calibrations determine the pointing of each detector at any time during the observations. The relation between pointing and time can depend on the following elements: \\begin{enumerate} \\item The ecliptic coordinates of the mean position of the spin axis during a TOP: $(\\lambda_z, \\beta_z)$ (see Fig.~\\ref{fig_ang}). This position and the position of the North ecliptic pole define a great circle, referred to as the `reference circle'. \\item For each TOP the nutation ellipse parameters $t_0$, $\\omega_{x,0}$, $\\omega_z$, $f_1$ and $f_2$ as defined in equation~(\\ref{eq_om_t0}) in Appendix~\\ref{app_wobble}. The first two of these determine the $x$-amplitude and phase of the nutation ellipse, the other three determine the period and the ratio between the $x$ and $y$ amplitudes. The nutation ellipse is described in the satellite coordinates of the inertial reference system, as defined in Appendix~\\ref{App_att}. When observed on the sky, the ellipse is itself rotating at the spin velocity of the satellite. The nutation period changes slowly due to variation in $f_1$ and $f_2$ resulting from depletion of consumables on-board the spacecraft. \\item The spin velocity $\\omega_z$ and its linear dependence on time during a TOP. The latter is again a parameter which changes very slowly over the mission. \\item For each TOP, the time of the first transit of the fiducial reference point (FRP) in the focal plane through the reference circle. A first approximation can be obtained from, for example, the star mapper data and the calibration of the star mapper\\footnote{A star mapper is a device recording the transits of bright stars for the purpose of spacecraft attitude control.} position with respect to the focal plane geometry. Corrections to the time of first transit are obtained during construction of the point-source catalogue. \\item The angle between the FRP and the mean spin axis position (the opening angle $\\alpha_0$, where the index `0' indicates the FRP rather than a specific detector `$d$'). This is a correction to the ground-based calibration value, that will change slowly with time due to inertia tensor changes. \\item The focal plane geometry, which describes, for an assumed orientation of the focal plane assembly, the coordinates of all detectors as projected on the sky, relative to the projection of the FRP in the focal plane. \\item The focal plane orientation correction, describing the difference between the actual and the assumed orientation of the focal plane as a function of time. \\end{enumerate} While items (i), (ii) and (iii) are probably best derived from the star mapper data, items (iv) to (vii) rely entirely on information contained in the science data, in particular on transits of bright point sources. Starting values for the focal plane geometry will be obtained from the instrument design specifications, but still need to be verified in flight. The calibration of these parameters relies on observed abscissae and intensities of point sources for different detectors. The geometric calibrations are crucial to the reliability of the final mission products, and inevitably require a number of iterations through the data reduction and calibration processes. As a part of these iterations the bright point-source catalogue (BPSC from here on) is produced\\footnote{This is not the same as the early release point source catalogue to be produced by {\\it Planck}, which is intended to be an early compilation of approximate coordinates and intensities of point sources detected by {\\it Planck} to be used in follow-up observations by different instruments}. The BPSC provides the best possible relative positions and intensities of all brighter point sources detected by {\\it Planck}, and is finally used to calibrate the overall positional reference system to the International Celestial Reference System (ICRS; Kovalevsky et al.\\ 1997). With the development of the BPSC, the calibration of the central beam profile can gradually be improved. In the first iteration through the data it is necessary to assume that all point sources pass through the centre of the beam. If, to a first approximation, the beam is represented as a two-dimensional Gaussian, then this assumption will do no harm: the width of the beam profile is in this case not a function of the ordinate of the transit. During subsequent iterations the ordinates of the transits can be calculated using the geometric calibrations and point-source coordinates contained in the catalogue. Ordinate information can also be incorporated in the beam profile calibrations. The construction of the catalogue ensures that in the final iteration a complete sample of point sources (up to a maximum ordinate depending on the beam profile) are taken into account in the ring analysis. This is essential for the use of such data in the sphere analysis \\cite{C1}. The iterations through the data reductions are essential due to the fact that many of the instrument calibration parameters have to be derived from the science data itself. This applies to both the response and the geometric calibrations. The precision with which these parameters can be determined is an essential part in the final data-quality verification. Poorly determined parameters will inevitably leave their mark on the final data products. The data analysis needs to identify any such parameters and their possible effect on the data as part of its preparation for the scientific exploration. \\section[]{The ring analysis} \\label{Sec_circ} The inputs to the ring analysis are the samples collected during a TOP for a single detector, supplemented by timing information (an absolute time for the first sample and a sampling length). It is assumed here that these intervals can be recognized from the satellite's housekeeping data. The reference position of the spin axis during a TOP is defined as the centre of the nutation ellipse actually described by the spin axis, and is derived from the star mapper data processing. Each detector $d$ describes a small reference circle with opening angle $\\alpha_d$ around the reference spin axis position. The use of rings as an intermediate step in the data analysis of a full-sky survey mission was first proposed by Lindegren (1979) for the Hipparcos mission. The idea was independently explored for CMB surveys by Delabrouille, G{\\'o}rski \\& Hivon (1998), who investigated the recovery of the CMB power spectrum directly from the ring data.. \\subsection[]{Signal representation} \\label{Sec_signal} The signal for the sky can be considered as consisting of three components: a continuous background, extended sources (which are not separated from the background) and point sources of various intensities. Translating this signal into the data observed on a given ring, a number of effects have to be taken into account: \\begin{itemize} \\item The satellite is rotating at a scan velocity which is not a fixed value and which may in addition change slightly during a TOP; \\item The sky signal is convolved with a beam profile and sampled; \\item The sky signal has a (not necessarily constant) background signal added to it, originating from the instrument itself; \\item The sky signal is represented on an arbitrary and possibly slightly variable readout scale. \\end{itemize} Some of these effects are removed in the ring reductions, others in the construction and calibration of the harmonic frequency maps \\cite{C1}. In the analysis of a TOP, two types of contributions are solved for: \\begin{enumerate} \\item The Fourier components, representing the structure in the underlying continuum of the microwave sky and extended sources: $C_0$, $C_n$ and $S_n$, with $n=1,\\dots,{n_{\\mathrm max}}$; \\item The intensities and abscissae for point sources (ordinates are assumed zero or obtained from the point-source catalogue): $I_k$ and $\\psi_k$, with $k=1,\\dots,s$, $s$ being the number of sources solved for. \\end{enumerate} In addition, the following properties of the data are resolved: \\begin{enumerate} \\item The detector response variations, such as changes in the background signal, $\\Delta b(t-t_{\\mathrm r})$, and drifts in the quantum efficiency, $\\Delta q(t-t_{\\mathrm r})$, where $t_{\\mathrm r}$ is an arbitrary reference time; \\item The covariance matrix of the measurement noise. This will be the noise on bin-averaged data rather than individual samples (see Section~\\ref{Sec_phase}). \\end{enumerate} The Fourier components for all TOPs corresponding to detectors in the same frequency channel are used to reconstruct the spherical multipoles of the frequency maps \\cite{C1}. This complements the ``ring-torus'' methods developed by Wandelt \\& Hansen (2001) for power spectrum estimation from the analysis of ring data from a special class of scans. The methods required for harmonic map-making also draw heavily on the harmonic-space convolution algorithms developed by Wandelt \\& G{\\'o}rski (2001) and Challinor et al.~(2000). The ring analysis is iterated with the construction of the BPSC (see Section~\\ref{Sec_point}) by means of the geometric calibrations mentioned in Section~\\ref{Sec_att}. The complete inclusion of identified point sources is also ensured through the BPSC. The success of this iteration determines the final pointing noise uncertainties and their contributions to the noise on the frequency maps. \\subsection[]{The time-to-scan-phase relation} The first requirement for the processing of the TOP data is an accurate determination of the scan velocity and its change with time over the interval covered by the TOP. The scan velocity $\\omega_z(t)$ determines the relation between time $t$ and a relative scan phase $\\psi$ for the individual sampling intervals: \\begin{equation} \\psi(t) = \\psi(t_{\\mathrm r}) + \\int^t_{t_{\\mathrm r}}\\omega_z(t){\\mathrm d}t. \\end{equation} The attitude simulations (Appendix~\\ref{App_att}) show that the scan velocity can to high accuracy be approximated by the scan velocity at reference time $t_{\\mathrm r}$, \\SV, and a constant scan-velocity drift, \\SVD, as \\begin{equation} \\omega_z(t) = \\omega_z(t_{\\mathrm r}) + (t-t_{\\mathrm r})\\dot{\\omega}_z, \\end{equation} giving the following relation for $\\psi(t)$: \\begin{equation} \\psi(t) = \\psi(t_{\\mathrm r}) + (t-t_{\\mathrm r})\\omega_z(t_{\\mathrm r}) + \\frac{1}{2}(t-t_{\\mathrm r})^2\\dot{\\omega}_z. \\label{equ_psi} \\end{equation} A further disturbance on the time to scan phase relation comes from the wobble of the spin axis, but this is very small as is shown in Appendix~\\ref{app_wobble}. The scan velocity and its variation will most probably be derived from the star mapper data. However, Section~\\ref{Sec_PSPar} shows how information on the scan velocity is, in principle, also present in the science data. There may occasionally be a discontinuity in the relation between time and scan velocity due to semi-discrete torques caused by the satellite being hit by a micrometeorite. Using equation~(\\ref{equ_psi}), phases can be assigned to each sample within the TOP. Sorting the data explicitly or implicitly (through a reference index) on $\\psi$ modulo $2\\pi$ produces the phase-ordered data (POD) for a TOP. The use of POD has also been explored by Wandelt \\& Hansen (2001) in the context of power spectrum estimation. The further processing of the POD involves the following steps, which are described in detail in the sections below: \\begin{enumerate} \\item Spike detection is performed on the POD, where spikes will show up more clearly due to the $\\sim 55$ times higher density of data points (Section~\\ref{Sec_spike}); \\item Phase binning of the data: This achieves a very significant compression of the data without significant loss of information (Section~\\ref{Sec_phase}); \\item Corrections for short-term detector response variations (Section~\\ref{Sec_ShResp}); \\item Point source identification: either from the data stream itself or from the point-source catalogue, providing abscissae, ordinates and intensities (Section~\\ref{Sec_PSPar}); \\item Signal fitting: solving for the Fourier components representing the continuum and corrections to some of the point-source parameters, and the noise spectrum (Section~\\ref{Sec_CirSol}); \\end{enumerate} Processes (i), (ii) and (iii) are applied only once to the data, while processes (iv) and (v) are part of the iteration with the BPSC construction. \\subsection[]{Spike detection and removal} \\label{Sec_spike} In the POD spikes will be much more conspicuous than in the TOD, as data points are compared with others at almost identical telescope pointings. Filters have to be developed that can reliably detect spikes as outlying points. All spikes are to be removed, while their times and intensities should be collected to allow for tests of their statistical properties (distribution over time and intensity). It may also be necessary to remove the sample(s) immediately following a spike in the TOD, and to compare observed response variations for detectors with the occurrence of spikes. Confusion between spikes and transient objects should be carefully avoided. Once the POD has been searched for, and cleaned from, spikes, it is ready for further analysis. \\subsection[]{Phase binning of the ring data} \\label{Sec_phase} A typical TOP for an HFI detector will contain around $7\\times 10^5$ samplings (200~Hz sampling frequency). The Fourier analysis is likely to require an $l_{\\mathrm max}$ value of around 2500, giving some 5000 unknowns (see Section~\\ref{Sec_resol}). The process of phase binning compresses the ring data without significantly modifying it (as would be the case when resampling), and as a result does not inflict a significant loss of information. The compression of the data, typically by a factor of 40 to 50, significantly reduces the processing time and data storage requirements. Relative to the original observations, the phase-binned data provides a much improved basis for the detection of point sources and detector response variations. The phase binning works for both point sources and the harmonic analysis of the background signal and is based on principles developed for, and used extensively in, the Hipparcos data analysis of the 1200~Hz modulated main detector signal (ESA, 1997; van Leeuwen, 1997). The basic relation between an observation (one sampling $O_i$ by a single detector) and its Fourier representation is given by \\begin{equation} O_i = C_0 + \\sum_{n=1}^{n_{\\mathrm max}}\\bigl[C_n\\cos n\\psi_i + S_n\\sin n\\psi_i\\bigr] + N_i, \\label{equ_fou} \\end{equation} where $\\psi_i$ is the phase of the observation and $N_i$ represents the instrument noise. We will use this representation to describe the signal with all bright point sources removed. The ring is divided into $m$ phase bins of equal length $2\\pi/m$. The phase at the centre of each bin is given by $\\Psi_j=2\\pi j/m$. Every observation is associated with a bin $j$, which turns equation~(\\ref{equ_fou}) into \\begin{eqnarray} O_i &=& C_0 + \\sum_{n=1}^{n_{\\mathrm max}}\\bigl[C_n\\cos n(\\Psi_j+{\\mathrm d} \\psi_{ij}) \\nonumber \\\\ &&+ S_n\\sin n(\\Psi_j+{\\mathrm d}\\psi_{ij})\\bigr] + N_i, \\label{equ_fou2} \\end{eqnarray} where ${\\mathrm d}\\psi_{ij}=\\psi_i-\\Psi_j$. This can be further expanded to \\begin{eqnarray} O_i & =& C_0 + \\sum_{n=1}^{n_{\\mathrm max}}\\biggl[C_n(\\cos n\\Psi_j\\cos n {\\mathrm d}\\psi_{ij}\\nonumber \\\\ &&-\\sin n\\Psi_j\\sin n{\\mathrm d}\\psi_{ij}) + S_n(\\sin n\\Psi_j\\cos n{\\mathrm d}\\psi_{ij}\\nonumber\\\\ && +\\cos n\\Psi_j\\sin n{\\mathrm d}\\psi_{ij})\\biggr] + N_i. \\label{equ_fou3} \\end{eqnarray} \\begin{figure*} \\centerline{\\psfig{file=MB792_fig2.eps,width=16truecm,angle=-90}} \\caption{Results for a binning experiment at the highest {\\it Planck} resolution: 5~arcmin, using 12500 bins. On the left are shown the actual values per phase bin as observed, on the right the histogram of the distribution of observed values, with an arrow indicating the expected mean value. Top graph: number of samples per bin; middle graph: phase dispersion per bin; bottom graph: mean phase offset per bin.} \\label{fig_bins} \\end{figure*} Phase binning provides weighted means per bin $j$ for the left- and right-hand sides of equation~(\\ref{equ_fou3}). The weights for all contributions in a bin can be taken equal. If $n_j$ denotes the number of samples in bin $j$, then the following relation is obtained for the mean signal in a bin: \\begin{eqnarray} {\\cal O}_j &=& \\frac{1}{n_j}\\sum_{i\\in j} O_i = C_0 + \\sum_{n=1}^{n_{\\mathrm max}}\\biggl[\\nonumber \\\\ &&C_n\\bigl( \\Sigma_{{\\mathrm c},j}\\cos n\\Psi_j - \\Sigma_{{\\mathrm s},j}\\sin n\\Psi_j\\bigr) \\nonumber \\\\ &&+S_n\\bigl( \\Sigma_{{\\mathrm s},j}\\cos n\\Psi_j + \\Sigma_{{\\mathrm c},j}\\sin n\\Psi_j\\bigr)\\biggr] + {\\cal N}_j, \\nonumber \\\\ \\label{equ_fou4} \\end{eqnarray} where $\\Sigma_{{\\mathrm s},j}$ and $\\Sigma_{{\\mathrm c},j}$ are defined below. By choosing the bin size sufficiently small, the sums over the bins can be estimated. Experience with the Hipparcos data \\cite{E1} showed that a coverage of a full cycle of the highest spatial frequency by 6 bins is sufficient, and allows for the following approximations to be made: \\begin{eqnarray} \\Sigma_{{\\mathrm s},j} &=&\\frac{1}{n_j}\\sum_{i\\in j}\\sin n{\\mathrm d}\\psi_{ij} \\approx \\frac{n}{n_j}\\sum_{i\\in j}{\\mathrm d}\\psi_{ij} = n\\langle{\\mathrm d}\\Psi_j\\rangle,\\nonumber \\\\ \\Sigma_{{\\mathrm c},j}&=&\\frac{1}{n_j}\\sum_{i\\in j}\\cos n{\\mathrm d}\\psi_{ij} \\nonumber \\\\ & \\approx & 1-\\frac{n^2}{2n_j}\\sum_{i\\in j}({\\mathrm d}\\psi_{ij})^2 = 1-n^2\\sigma^2_{\\Psi_j}. \\label{equ_fou5} \\end{eqnarray} Depending on the amplitude of the nutation ellipse relative to the beam width, it may be necessary to fit the data contained in a phase bin as a second-order function of the offset from the reference ring position, primarily to avoid an intensity bias and noise increase due to point sources. The phase binning thus requires the accumulation of four items per bin: the mean count ${\\cal O}_j$, the total number of samples $n_j$, the average phase correction $\\langle{\\mathrm d}\\Psi_j\\rangle=\\sum{\\mathrm d}\\psi_{ij}/n_j$ and the phase dispersion $\\sigma_{\\Psi_j}=\\sqrt{\\sum({\\mathrm d} \\psi_{ij})^2/2n_j}$. Values for $\\langle{\\mathrm d}\\Psi\\rangle$ and $\\sigma_\\Psi$ are shown in Fig.~\\ref{fig_bins} for a simulation using a sampling frequency of 200~Hz, ${n_{\\mathrm max}}=2050$, and 12500 bins. The offset from the nominal scan velocity was 1.243~arcsec~${\\mathrm s}^{-1}$, and the scan velocity drift was 0.009~arcsec~${\\mathrm s}^{-2}$. An average of 57 samples per bin is expected under these conditions. The expected value for $\\langle{\\mathrm d}\\Psi\\rangle$ is zero, and $\\langle\\sigma_\\Psi\\rangle = (\\pi/m)/\\sqrt{6}$ (where $m$ is the total number of bins), which equals $21\\farcs 6$, equivalent to 5.8~per cent of the beam dispersion, or 5~per cent of the bin width in this example. The phase-binned data ${\\cal O}_j$ and $\\langle{\\mathrm d}\\Psi_j\\rangle$, $\\sigma_{\\Psi_j}$ and the correlation matrix of the ${\\cal N}_j$ enter our model of the response through equations~(\\ref{equ_fou4}) and (\\ref{equ_fou5}) in a generalised least squares solution for the $\\{C_n,S_n\\}$. In the evaluation of equation~(\\ref{equ_fou5}) we ignored the fourth order term, which has an expected value of $(n\\pi/m)^4/120$. With $m=6n$ for the highest $n$ value used, this amounts to a maximum correction of $6\\times 10^{-4}$ to $\\Sigma_{{\\mathrm c},j}$. This figure should be compared with the expected amplitudes for the highest frequency harmonics (which will be severely depressed by the beam convolution), relative to the noise on the measurement ${\\cal O}_j$. For the lower frequencies the contribution of the fourth order term is very much smaller still (e.g.~$6\\times 10^{-12}$ for $n=25$). In fact, for most of the lower frequencies the approximations of $\\Sigma_{\\mathrm s}\\approx 0$ and $\\Sigma_{\\mathrm c}\\approx 1$ can be used. Thus, phase binning brings down the number of observations used for the ring analysis by a factor $\\sim 57$, without any significant loss of information. The data storage requirements are as a result significantly reduced. All trigonometric coefficients in the analysis are fixed in the binned solution, and can be pre-calculated. The phase-binned data for each ring can be kept as intermediate data products, as the contents do not change in the iterative processes. \\subsection[]{Short-term response variations} \\label{Sec_ShResp} As stated before (Section~\\ref{Sec_signal}), the short-term response variations can be derived from the systematic differences between the individual sample counts and the mean count of the phase bin to which the sample has been assigned. Correlated noise between neighbouring samples will transfer to similar noise between neighbouring bins. Two types of corrections can be made: \\begin{itemize} \\item A background variation correction $\\Delta b(t-t_{\\rm r})$, which can represent effects like changes in the radiation received from the instrument and the spacecraft, and which acts like a varying constant offset; \\item A quantum efficiency variation correction $\\Delta q(t-t_{\\rm r})$, which will produce residuals scaled by the mean bin count. \\end{itemize} For a sample $O_i$ in phase bin $j$ with mean count ${\\cal O}_j$, this means \\begin{equation} \\Delta O_{ij} \\equiv O_i - {\\cal O}_j = \\Delta b(t-t_{\\rm r}) + \\Delta q(t-t_{\\rm r}){\\cal O}_j + N_i. \\end{equation} The functional representation for these two corrections has to be decided upon early in the mission, on the basis of the features shown in the data. A fit with a low-order spline function will in most cases take care of any real variations. At this stage it is unnecessary, and potentially even damaging, to reduce the detector responses to an absolute scale, in particular if those calibrations would include a dependence on the phase angle $\\psi$, for which the reference phase has not been accurately determined yet. The observed variations should be compared with temperature records for the spacecraft and payload and the occurrences of spikes, and correlations may be used in the removal of any observed variations. As a result of this calibration, the data collected during each TOP will be free from short-term detector variations. \\subsection[]{Corrections for the satellite's motion} \\label{sec_aberr} The {\\it Planck} satellite will be positioned in a Lissajous orbit around the L2 Lagrangian point of the Sun-Earth system. It will therefore describe an almost circular orbit about the Sun with a one year period and a radius of $\\sim 1.01$~AU. The period for the Lissajous orbit relative to L2 is around 179 days, and will have an amplitude of around $10^5$~km. Thus, the orbital velocity of the satellite is dominated by its motion around the Sun and will be approximately 30.3~km~s$^{-1}$. This causes fractional spectral shifts of $\\Delta\\lambda/\\lambda\\approx 10^{-4}$, which is equivalent to 9 per~cent of the dipole signal in the CMB radiation. The {\\it Planck} mission aims at detecting much smaller anisotropies in the CMB, and these effects are therefore a significant distortion of the signal. The effect will be opposite in the two half-year surveys, and will be most noticeable near the ecliptic plane. The data can be corrected for this effect iteratively with the production of the frequency maps. The frequency maps can be prepared to relatively low values of $l_{\\mathrm max}\\approx500$ to produce all-sky spectral index maps. The velocity vector of the satellite together with estimated maps of the spectral gradient can then provide corrections to the observed intensities for each ring: \\begin{equation} \\Delta I \\approx \\lambda\\frac{\\partial I}{\\partial \\lambda}\\frac{v}{c}\\cos\\theta, \\end{equation} where $\\theta$ is the angle between the velocity vector of the observer (which has magnitude $v$) and the observation direction, $c$ is the speed of light in vacuum, $\\partial I/\\partial \\lambda$ is the local spectral gradient, and $\\Delta I$ the local intensity correction. This effect then has to be integrated over the spectral response of the beam profile to correct the actual observed signal. The positional effect, generally referred to as aberration, is, to first order in $v/c$, given by \\begin{equation} \\sin\\Delta\\theta = (v/c)\\sin\\theta, \\end{equation} where $\\Delta\\theta$ is the difference between the propagation direction of the radiation in a stationary reference frame and the actual moving reference frame. For the {\\it Planck} observations $\\Delta\\theta$ has a maximum of 20~arcsec (for the part of the scan nearest to either of the ecliptic poles) and can in principle be taken into account when assigning scan phases to the individual samples. This is likely to be relevant for the HFI detectors, for which the maximum correction is comparable with the abscissa accuracies for bright point sources, and the angular scale of the highest harmonics used in the signal analysis. Ignoring the correction will result in information leakage between neighbouring frequencies in the harmonic analysis, and a significant positional noise contribution for point sources near to the ecliptic poles. \\subsection[]{Resolution requirements} \\label{Sec_resol} An important aspect of the realization of the harmonic data model is the value of ${n_{\\mathrm max}}$ that should be applied to the data analysis. This value is largely determined by the beam width, the scan density, and the way in which point sources are dealt with in the analysis. Good estimates for ${n_{\\mathrm max}}$ can be obtained by assuming a circular-symmetric Gaussian beam profile. The power spectrum of a point source convolved with a Gaussian beam is shown in Fig.~\\ref{fig_point}. It is clear that in order to represent the point sources adequately as components in the harmonic analysis relatively high values of $n_{\\mathrm max}$ are required. An alternative approach is to treat identifiable point sources as separate components in the ring analysis. This should be feasible if no features are expected in the background signal, which are sharp compared to the beam profile. \\begin{figure} \\centerline{\\psfig{file=MB792_fig3.eps,width=8truecm}} \\caption{Bottom graph: The power spectrum of a point source (with sampled peak intensity equal to 1) given a Gaussian beam with \\textsc{fwhm} of 5~arcmin ($\\sigma_b=2.12$~arcmin) and a sample width of 1.8~arcmin. Middle graph: the restored peak intensity of a point source if fitted with harmonic components up to the indicated $n_{\\mathrm max}$ value. Top graph: the restored point source for $n_{\\mathrm max}=2500$ (thick), 4000 (intermediate) and 6000 (thin) (also distinguished by increasing peak height and decreasing side-lobe amplitudes).} \\label{fig_point} \\end{figure} A circular symmetric Gaussian beam with a dispersion $\\sigma_b$ (in radians) will reproduce the higher harmonics in the background signal with decreased amplitudes. If the beam is represented by \\begin{equation} R(\\psi) = \\exp[-\\psi^2/2\\sigma_b^2], \\end{equation} then the decrease in the amplitude is given by \\begin{equation} {a_{lm}}' = a_{lm}\\exp[-(l\\sigma_b)^2/2], \\end{equation} where ${a_{l}}$ is the actual and ${a_{lm}}'$ the observed amplitude (see also Challinor et al., 2002). Provided that point sources are treated as separate objects, the suppression of higher harmonics determines the value for $n_{\\mathrm max}$ in the ring solution. For $\\alpha_d\\approx 85\\degr$, simulations (to be detailed in a future paper) suggest a value of $n_{\\mathrm max}\\approx 4400~{\\mathrm arcmin} /\\sigma_b$, where $\\sigma_b$ is expressed in arcmin. Thus, for $\\sigma_b=2.12$~arcmin (\\textsc{fwhm} of 5~arcmin) we find $n_{\\mathrm max}\\approx 2100$. At such high $n_{\\mathrm max}$ values the matrices involved in the transformations are very large, and the covariance matrix of the spherical multipole solution will contain around $10^{13}$ elements. Fortunately, simulations, which will be detailed in a future paper, have shown that for many plausible scanning strategies, the covariance matrix of this solution will be very sparse (some analytic approximations to the covariance matrices for simple scan strategies are also derived in Challinor et al., 2001). \\subsection[]{The point-source parameters} \\label{Sec_PSPar} The main tasks of the point source processing are the following: \\begin{itemize} \\item To identify from either the data stream or the BPSC the point sources present in the ring data; the source of information will depend on the iteration stage of the data reduction process; \\item To supplement this list with solar system objects that may have been observed; \\item To produce for each point source preliminary estimates of the intensity $I_s$, the abscissa $\\psi_\\tau$, and if obtained from a catalogue or ephemerides (solar system objects), the ordinate $\\upsilon_\\tau$; \\item To obtain as part of the reduction of the ring data the intensities $I_s$ of all identified sources, and abscissae $\\psi_\\tau$ for sources with a sufficiently high signal-to-noise ratio; \\item To collect the abscissa data for (re-)building the BPSC; \\item To collect the profiles and fitting parameters for the brightest sources as a contribution to the beam profile calibration. \\end{itemize} This process is clearly iterative, improving at each step the quality of the BPSC, the beam profile and the reliability of the segregation of the point sources from the data. \\subsubsection[]{Identification of point sources} \\label{Sec_IdentPS} The mechanism of point source identification will depend on the iteration phase of the data reductions: in the first instance point sources will mainly be identified from the data stream itself, while in subsequent iterations the identification will rely more on the BPSC. It is expected that, at least for the high-frequency detectors, the power will be dominated at high $l$ values by point sources, which should allow for the development of a reliable filter for the detection of brighter point sources. Filters for the recognition of point sources in a one-dimensional data stream were developed and successfully applied for the Hipparcos and Tycho data reductions \\cite{L1}. Simplified versions of two-dimensional filters under development for recognition of point sources on maps (see for example Cayon et al., 2000, Sanz et al., 2001) could also be considered. The detection of point sources from the data is significantly enhanced by the phase binning of the data, increasing the signal-to-noise ratio by almost a factor 7 for the parameters used in Section~\\ref{Sec_phase}. On the first pass through the data all point sources are assumed to transit through the centre of the beam. This is not problematic if the beam is approximately circularly symmetric and Gaussian. The deviation of the actual beam profile from this assumption will cause a slight error in the first reconstruction of the beam. This error can be reduced once ordinate information on point source transits becomes available too. All detected point sources will be used to build the first and subsequent versions of the BPSC: the evolving catalogue used in iterations to predict and consistently identify point sources. In these iterations information on the ordinate of the source at the time of the transit should also be incorporated. The consistent inclusion of point sources in the ring analysis is an essential requirement for the further analysis of the underlying continuum. The BPSC also plays a crucial role in the geometric calibrations of the instrument. \\subsubsection[]{The measured and binned point-source signal} \\label{Sec_MeasPSSgn} \\begin{figure} \\centerline{\\psfig{file=MB792_fig4.eps,width=7truecm}} \\caption{The point-source position (PS) relative to the phase binning and the mean scan circle.} \\label{fig_PS} \\end{figure} The sampled responses $O_i$ for a point source of intensity $I_s$, passing through the beam of a detector, is a function of the scan phase $\\psi_i$ (at the midpoint of the sampling interval) for sample $i$ and the abscissa $\\psi_\\tau$ and ordinate $\\upsilon_\\tau$ of the point source (see Fig.~\\ref{fig_PS}): \\begin{equation} O_i=I_s\\int^{\\psi_i+\\Delta\\psi}_{\\psi_i-\\Delta\\psi} R(\\psi_\\tau-\\psi,\\upsilon_\\tau){\\mathrm d}\\psi, \\end{equation} where the integral represents the sampling interval, and $R(\\psi_\\tau-\\psi,\\upsilon_\\tau)$ is the normalized beam profile for a detector as a function of the offset from the centre of the beam ($\\psi_\\tau-\\psi$ along the scan direction and $\\upsilon_\\tau$ perpendicular to the scan direction). \\begin{figure} \\centerline{\\psfig{file=MB792_fig5.eps,width=8truecm}} \\caption{The beam response $R(\\psi,0)$ (thin line) and the convolved beam response $S(\\psi,0)$ (thick line) profiles for a Gaussian beam with \\textsc{fwhm}$=5$~arcmin. The increase in beam width as a result of the samplingis of the order of 20~arcsec. Top graph: in linear response scale; bottom graph: in logarithmic response scale.} \\label{fig_beam} \\end{figure} The effective beam profile $S(\\psi,\\upsilon)$ is defined as the actual beam profile $R(\\psi,\\upsilon)$ convolved with the sampling interval, as illustrated in Fig.~\\ref{fig_beam} [see also equation~(\\ref{equ_fou5})] \\begin{equation} S(\\psi,\\upsilon) = \\int^{\\psi+\\Delta\\psi}_{\\psi-\\Delta\\psi} R(\\psi',\\upsilon) {\\mathrm d}\\psi', \\end{equation} where the relevant coefficients are defined in Section~\\ref{Sec_phase}. When phase binning is applied to the point source contributions, the response in phase bin $j$ is given by \\begin{eqnarray} {\\cal O}_j &=& I_s\\bigl[S(\\psi_\\tau-\\Psi_j,\\upsilon_\\tau) - \\langle{\\mathrm d}\\Psi_j\\rangle S^\\prime(\\psi_\\tau-\\Psi_j,\\upsilon_\\tau) \\nonumber \\\\ && + \\sigma_{\\Psi_j}^2 S^{\\prime\\prime}(\\psi_\\tau-\\Psi_j,\\upsilon_\\tau)\\bigr] + {\\cal B}_j + {\\cal N}_j, \\label{eq:ps_bin} \\end{eqnarray} where ${\\cal B}_j$ represents the background signal. In the same way as was found for the harmonic multipoles, the main contribution comes from the second derivative, which produces an effective broadening of the beam. However, if we use 12~500 bins for a beam with \\textsc{fwhm}=5~arcmin ($\\sigma_b=2.12~{\\mathrm arcmin}$), the additional dispersion is $\\sigma_{\\Psi_j}\\approx 21\\farcs 6$ and the effective \\textsc{fwhm} of the beam is increased as a result of the phase binning by no more than one per~cent. \\subsubsection[]{Fitting parameters for the point-source signal} Two parameters require determination in the signal fit: the transit phase $\\psi_\\tau$ and intensity $I_s$ of the point source. For both, preliminary estimates are required which are then adjusted in the solution. Equation~\\ref{eq:ps_bin} can be linearized in the intensity correction $\\delta I_s$ and the transit correction $\\delta\\psi_\\tau$ so that \\begin{eqnarray} {\\cal O}_j &=& \\delta I_s S(\\psi_\\tau-\\Psi_j,\\upsilon_\\tau) + \\delta\\psi_\\tau I_s S^\\prime(\\psi_\\tau-\\Psi_j,\\upsilon_\\tau) \\nonumber \\\\ && + I_sS(\\psi_\\tau-\\Psi_j,\\upsilon_\\tau) + {\\cal B}_j + {\\cal N}_j. \\label{equ_point} \\end{eqnarray} This equation enters in the generalised least squares solution for the ring data (see Section~\\ref{Sec_CirSol}). The intensity correction term in Eq.~\\ref{equ_point} is scaled by the response function. As a result of this, the noise on the lower intensities (the wings of the response function) will tend to affect the determination of $\\delta I_s$ more than the better determined higher intensities (core of the response function). This is in particular the case when the noise on the signal depends on its intensity. Therefore, only a few of the central phase bins should be used for solving for the transit parameters of a given source. It may be necessary to iterate the ring solution to properly separate the point source and continuum contributions. Assuming we use the central 5 phase bins for determining the fitting parameters for each point source, estimates can be obtained for the expected precisions. For the intensity error a value of 0.8 times the noise level on the binned samples is expected, equivalent to 0.12 times the noise level on individual samples. For the abscissa error we estimate a value of 1.7 arcmin divided by the signal-to-noise ratio of the point source in the phase-binned data. \\subsubsection[]{Transient sources} Any object moving by a significant fraction of the beam width over an interval of one hour will have to be treated as a transient source. For the highest frequency channels of {\\it Planck} this translates into a movement of approximately 2~arcmin per hour and above. The fastest objects in the Earth's neighbourhood will be traveling at $\\sim 10^5$~km~hr$^{-1}$. The limit of 2 arcmin then translates into a horizon of around 1.2 AU, which will include a small fraction of asteroids and the occasional comet. For sources that are not variable on a time-scale of less than one hour (which may not be true due to rotation of the objects), systematic differences between the actual measurements (before phase binning) and the reference profile can be expressed as corrections to the effective scan velocity \\SV: \\begin{eqnarray} \\Delta O_i &=& O_i - B_i - I_sS(\\psi_\\tau-\\psi_i,\\upsilon) \\nonumber \\\\ &\\approx& I_s\\frac{\\partial S}{\\partial\\psi} \\Delta(\\psi_i-\\psi_\\tau) + N_i \\nonumber \\\\ &=& (t-t_{\\mathrm r})I_s\\frac{\\partial S}{\\partial\\psi}{\\mathrm d}\\omega_z + N_i, \\label{eq_phas2} \\end{eqnarray} which shows that most of the information on \\SV is contained in those sections of the signal which have the steepest gradient as a function of $\\psi$. The same kind of information can in principle also be used to determine the scan velocity from the science data using bright-point-source measurements, although it would be by far preferred to derive this information from the star mapper data. \\subsubsection[]{Beam profile calibration} The beam profile calibration uses the accumulated transit data of bright point sources (which may include transits of solar system objects). During the first processing of the data no reliable ordinate information is available (catalogue positions are still poor or non-existent, and opening angles still need to be calibrated), and every source is assumed to go through the centre of the beam. During later stages of the data reductions, when the point-source catalogue and the geometric calibrations are improving, the ordinate information can be incorporated too. The principle of the beam profile calibration follows techniques used in the Hipparcos data reductions for the single slit response functions of the star mapper slits (ESA, 1997; van Leeuwen, 1997). The response function is sampled at a frequency 4 to 8 times higher than the sampling frequency of the data. Using the abscissae and intensities of identified point sources, the observed residual samplings ${\\cal O}_j-{\\cal B}_j$ (original signal ${\\cal O}_j$ minus the background ${\\cal B}_j$ as estimated from the current harmonic fit to the continuum) in phase bins close to the point source transit phases $\\{\\psi_\\tau\\}$ are binned according to their separation from the transit phase $\\{\\psi_\\tau\\}$. Also accumulated in bins are the derived transit peak intensities $I_s$. Thus, we find for the beam-profile bin $k$ the mean normalised response \\begin{equation} {\\cal S}(k) = \\frac{\\sum \\bigl[{\\cal O}_j(k)-{\\cal B}_j(k)\\bigr]}{\\sum I_s(k)}, \\label{equ_beam} \\end{equation} where the sums are over the point sources and phase bins contributing to the beam-profile bin $k$. Noise correlations between neighbouring bins in the phase-sampled data will also affect the accumulation of equation~(\\ref{equ_beam}), in that pairs of bins in the accumulation can contain partially correlated noise, which needs to be taken into account when fitting a response curve. The values ${\\cal S}(k)$ can be fitted with a cubic spline to provide a smooth beam profile with a continuous derivative, which then provides an estimate for the sampled beam profile $S(\\psi,0)$. When, after the construction of the first point-source catalogue, ordinates become available for the point-source transits, the beam profile can be resolved in two dimensions. The reconstructed beam profile obtained this way is the sampled, effective profile, and not the actual profile, which would apply to a stationary detector. Transits of planets such as Jupiter may also be useful for the beam profile calibration, though could be problematic: the detector response to very high intensities will not be linear, and the spectral gradient for the planets is likely to be quite different from that of the average microwave point source. The beam profiles will vary with frequency (see for example Challinor et al., 2001), making it still more complicated to incorporate the profiles measured from the planets. \\subsection[]{The ring solution} \\label{Sec_CirSol} After binning the data, calibrating and removing short-term detector responses, and identifying the point sources, the TOP is ready for reduction. This part of the reductions consists of a generalised least squares solution for the ring harmonics in the underlying continuum (equations~(\\ref{equ_fou3})--(\\ref{equ_fou5})) and simultaneous solutions for the abscissa and intensity corrections for all identified point sources (equation~(\\ref{equ_point})). The observation vector $\\bmath{z}$ has as its components the mean response in each phase bin. The observations are related to the vector $\\bmath{x}$ containing the amplitudes of the circle harmonics, $\\{C_n, S_n\\}$, and the corrections to the point source parameters, $\\delta I_s$ and $\\delta \\psi_\\tau$, and the noise vector $\\bnu$, whose components are the noise in each phase bin, ${\\mathcal{N}}_j$, via the linear equation \\begin{equation} {\\bmath z} = {\\mathbfss A}{\\bmath x} + {\\bnu}. \\label{eq_obs} \\end{equation} The matrix $\\mathbfss{A}$ depends on the locations of the phase bins, $\\Psi_j$, the phase corrections and dispersions, $\\langle {\\mathrm{d}} \\Psi_j \\rangle$ and $\\sigma_{\\Psi_j}$, the current estimate of the point source intensities and positions, $I_s$ and $\\psi_\\tau$ (and ordinate information as this becomes available), and the current estimate of the beam profile, $S(\\psi,v)$. Equation~(\\ref{eq_obs}) is the matrix form of equations~(\\ref{equ_fou4}) and (\\ref{equ_point}). The minimum-variance estimate of $\\bmath{x}$ is \\begin{equation} \\hat{{\\bmath{x}}} = ({\\mathbfss{A}}^T {\\mathbfss{N}}^{-1} {\\mathbfss{A}})^{-1} {\\mathbfss{A}}^T {\\mathbfss{N}}^{-1} {\\bmath{z}}, \\end{equation} with errors \\begin{equation} \\langle (\\hat{{\\bmath{x}}} - {\\bmath{x}})(\\hat{{\\bmath{x}}} - {\\bmath{x}})^T \\rangle = ({\\mathbfss{A}}^T {\\mathbfss{N}}^{-1} {\\mathbfss{A}})^{-1}, \\end{equation} where ${\\mathbfss{N}} \\equiv \\langle \\bnu \\bnu^T \\rangle$ is the (phase-binned) noise covariance matrix. Assuming the instrument noise is a stationary, random process with correlation function $C(t)$ in the time domain, we can obtain the noise contribution to the mean response in a given phase bin by integrating over the time periods corresponding to those observations falling in that bin. The covariance matrix $[{\\mathbfss{N}}]_{jj'}\\equiv \\langle {\\mathcal{N}}_j {\\mathcal{N}}_{j'}\\rangle$ then takes the form of a convolution: \\begin{equation} [{\\mathbfss{N}}]_{jj'} = \\left( \\frac{m}{N_{\\rmn{s}}} \\right)^2 \\int_{x_-}^{x_+} \\Lambda(x) C\\{T_{\\rmn{s}}[x + (j-j')/m]\\}\\, {\\rmn{d}} x, \\end{equation} where $N_{\\rmn{s}}$ is the number of times the ring is scanned in one TOP, $T_{\\rmn{s}}$ is the average spin period in that TOP, $m$ is the number of phase bins, and the integration limits $x_\\pm \\equiv -(N_{\\rmn{s}}-1)\\pm 1/m$. The function $\\Lambda(x)$ is given by \\begin{eqnarray} \\Lambda(x) &=& \\sum_{n=-(N_{\\rmn{s}}-1)}^{N_{\\rmn{s}}-1}\\Bigl[ \\Theta(1/m - |x-n|)\\nonumber \\\\ && \\qquad\\times(1/m - |x-n|)(N_s-|n|)\\Bigr], \\end{eqnarray} where $\\Theta(x)$ is the Heaviside unit step function, and arises from the effects of phase binning and repeatedly scanning the ring. For white noise the ${\\mathcal{N}}_j$ are uncorrelated, but in the presence of a significant low frequency component correlations will arise. Typically, instruments are designed with the goal of restricting coloured noise to frequencies below the spin frequency, in which case the correlation length of the noise exceeds the spin period. (For \\emph{Planck} HFI, the nominal knee frequency at which the character of the noise changes from $1/f$ to white is less than $0.06\\, \\rmn{rad}~\\rmn{s}^{-1}$, while the spin frequency is $0.10\\,\\rmn{rad}~\\rmn{s}^{-1}$.) As the correlation length becomes large compared to the spin period, the noise covariance matrix approaches $[{\\mathbfss{N}}]_{jj'} = \\chi^2_{\\rmn{c}} + \\chi^2_{\\rmn{u}} \\delta_{jj'}$, corresponding to a fully correlated offset in every bin with r.m.s.\\ $\\chi_{\\rmn{c}}$, and uncorrelated noise with r.m.s.\\ $\\chi_{\\rmn{u}}$. In practice, the noise power spectrum will have to be estimated from the data directly rather than relying on simple parameterised forms like that give above. If the correlation length exceeds $N_{\\rmn{s}} T_{\\rmn{s}}$ the offsets on different rings will generally also be correlated, so an optimal analysis would require reducing several rings simultaneously. \\begin{figure} \\centerline{\\psfig{file=MB792_fig6.eps,width=8truecm}} \\caption{A histogram of the normalised, absolute values of the off-diagonal elements in the square root of the covariance matrix for a typical bin-sampled ring solution, using $l_{\\mathrm max}=512$. The matrix can effectively be regarded as diagonal.} \\label{fig_hist} \\end{figure} Under the assumption of Gaussianity, the relevant components of the covariance matrix $\\langle (\\hat{{\\bmath{x}}} - {\\bmath{x}})(\\hat{{\\bmath{x}}} - {\\bmath{x}})^T \\rangle$ determine the errors on the ring harmonics (marginalised over the point source corrections). We have conducted several simulations to investigate the effects of variations in the spin velocity, and the presence of point sources, on the covariance matrix for the errors on the ring harmonics. To isolate the effects of spin velocity and point sources we have only considered white instrument noise. For our simulations we adopted a maximum value of $n$ equal to $512$. The covariance matrix for the errors on the ring harmonics is found to be very close to diagonal if only relatively faint point sources are present. In the presence of a very bright point source (like Jupiter), which will create a very low-weight gap in the ring data, there is a minor effect on the diagonal structure of the covariance matrix, the effects being largest in low frequency detectors. The normalised off-diagonal elements in the (Cholesky) square root of a typical covariance matrix are accumulated in a histogram in Fig.~\\ref{fig_hist}. The results show that the correlations between the different ring harmonics are at a level of 0.2~per cent or less. These levels decrease still further with increased resolution. The inclusion of a bright point source would affect the distribution, but correlations would still be below the 1~per cent level. If we further include stationary instrument noise, we can expect the errors on the ring harmonics to still be uncorrelated between Fourier modes for large $N_{\\rmn{s}}$. If coloured noise is confined to frequencies well below the spin frequency, correlations between rings will only affect the $n=0$ modes~\\cite{C1}. \\section[]{The bright-point-source catalogue} \\label{Sec_point} The bright-point-source catalogue (BPSC) is both a major product of the {\\it Planck} mission, and an important calibration tool. The catalogue is constructed iteratively from the abscissae and intensities of the point sources obtained in the current ring reductions. Corrections to the BPSC will ultimately provide positions, intensities (at different frequencies), and variability information for all detected sources. As was explained in Section~\\ref{Sec_calibr}, the spin axis position (including nutation movements) and spin rate are most likely derived from the star mapper data, all other geometric calibrations depend at least to some extent on the science data. \\subsection[]{The positional sphere reconstruction} The construction of the BPSC follows a simplified version of the Hipparcos sphere reconstruction process \\cite{E2}. This process requires a set of a priori positions for the point sources, to which corrections are determined. These positions are also required in order to identify consistently point-source transits on different rings with objects on the sky. Initial positions at a precision of twice the beam size (10~arcmin for {\\it Planck}) will be sufficient for this purpose. Using the ground-based specifications of the focal plane geometry and inertia tensor, first estimates are obtained for the opening angle and focal plane rotation corrections, and the focal plane geometry. These estimates can be further refined through the use of planetary transits. From this information we obtain preliminary details of the alignment of the different detectors during a scan: phase offsets and effective opening angles. The transits recorded in scans for the different detectors can now be transferred to a common sky, where each transit is represented by an error ellipse, strongly elongated in a direction perpendicular to the scan direction. A simple algorithm is then required to cross identify the different transits, and to produce from the accumulated information the necessary initial positions. Once a catalogue of initial positions has been obtained, and transits have been identified with sources in the catalogue, corrections to the assumed positions and to the reference phases of the rings can be determined. Given the ecliptic coordinates $(\\lambda_z,\\beta_z)$ of the spin axis, and a position for a point source $i$ at $(\\lambda_i,\\beta_i)$, the expected abscissa and ordinate of the point source with respect to the ring can be derived. The coordinates $(\\psi_i,\\zeta_i)$ in the system with the spin axis at the pole (where $\\psi_i$ is measured from the reference circle as defined in Fig.~\\ref{fig_ang}, and $\\zeta_i$ is the latitude of the source relative to the great circle associated with the spin-axis position) are obtained from \\begin{equation} \\left[\\matrix{\\cos\\psi_i\\cos\\zeta_i\\cr\\sin\\psi_i\\cos\\zeta_i\\cr\\sin\\zeta_i}\\right] = {\\mathbfss R}_2\\bigl(\\beta_z-\\frac{\\pi}{2}\\bigr){\\mathbfss R}_3(-\\lambda_z) \\left[\\matrix{\\cos\\lambda_i\\cos\\beta_i\\cr\\sin\\lambda_i\\cos\\beta_i\\cr \\sin\\beta_i}\\right], \\label{eq_coor_tr} \\end{equation} where ${\\mathbfss R}_i(\\phi)$ is the matrix representing a right-handed rotation around axis $i$ through angle $\\phi$. Equation~(\\ref{eq_coor_tr}) can be used to derive the relation between corrections to the source position $(\\lambda_i,\\beta_i)$ and the resulting changes to the scan phase $\\psi$ (obtained from the transit abscissae) and ordinate $\\upsilon$ (reflected in the transit intensities). We define $\\zeta=\\frac{\\pi}{2}-\\alpha+\\upsilon$, where $\\upsilon$ is the ordinate relative to the actual ring, and $\\alpha$ the opening angle. The relations between the offset values for $(\\Delta\\psi_i,\\Delta\\upsilon_i)$ and $(\\Delta\\lambda_i,\\Delta\\beta_i)$ are then found to be \\begin{eqnarray} \\Delta\\psi_i\\cos^2\\zeta_i &=& f_1\\Delta\\lambda_i\\cos\\beta_i+f_2\\Delta\\beta_i, \\nonumber \\\\ \\Delta\\upsilon_i\\cos\\zeta_i &=& -f_2\\Delta\\lambda_i\\cos\\beta_i+ f_1\\Delta\\beta_i, \\end{eqnarray} where $f_1$ and $f_2$ are defined as \\begin{eqnarray} f_1 &=& \\sin\\beta_z\\cos\\beta_i-\\cos\\beta_z\\sin\\beta_i\\cos(\\lambda_i-\\lambda_z), \\nonumber \\\\ f_2 &=& \\cos\\beta_z\\sin(\\lambda_i-\\lambda_z). \\end{eqnarray} For a scanning strategy where the spin axis remains in or close to the ecliptic plane (and $\\lambda_i-\\lambda_z\\approx\\pm\\pi/2$), $f_1\\approx 0$ and $|f_2|\\approx 1$ for most sources, the exception being sources situated close to either of the ecliptic poles. The result is that very little information on the ecliptic longitudes can be extracted from the measured abscissae, except for images at high ecliptic latitudes. Some information on the longitudes can, however, be extracted using the measured intensities, which depend on $\\upsilon$ through the beam profile. This method will inevitably fail when the source is variable. The positional sphere solution consists of a least squares fit of the differences between the observed and predicted abscissae, $\\Delta\\psi_{i,j}$ (for point source $i$ on ring $j$) to corrections $(\\Delta\\lambda_i,\\Delta\\beta_i)$ for each point source. In addition, there is a correction $\\Delta\\psi_j$ to the assumed reference phase for each ring. In the solution a boundary condition needs to be included which states that the average correction $\\sum\\Delta\\psi_j=0$, else a singularity may occur. When using different detectors, the phase offsets between the detectors have to remain effectively fixed: the only changes can come from rotation variations of the field of view, and focal length variations of the telescope, and both effects are unlikely to be significant. To estimate the number of observations and parameters, we assume 10~000 point sources on the sky bright enough to be detected on single rings (Barreiro, private communication), and 10~000 rings per detector over the {\\it Planck} mission. With the beam's \\textsc{fwhm}$=5$~arcmin, the average number of point sources detected per ring is 6 passing within $\\pm\\sigma_b$, and 6 more passing within $\\pm2\\sigma_b$ from the mean position of the ring. The total number of transits recorded is then $\\sim10^5$ per detector, giving $\\sim 10$ observations per source per detector. The total number of parameters to be estimated is $\\sim3\\times10^4$, and this does not increase when information from more detectors is used (except when some point sources are not visible for all detectors used). Increasing the number of detectors increases the number of observations and will increase the rigidity of the solution. The relative positions of point sources obtained this way are rigid, but absolute positions are only determined up to an overall rotation, and so require linking to the International Celestial Reference Frame \\cite{K1}. This can be accomplished through cross identification with radio and optical counterparts, thus determining a global rotation to be applied to the entire catalogue. This may not be important for the statistical properties of the CMB component in the background, but is relevant for relating sources and the dust map to other observations. \\subsection[]{Geometrical calibrations} \\label{Sec_geom} In Section~\\ref{Sec_calibr} we mentioned four geometric calibrations which rely on the point-source data collected by the satellite: \\begin{enumerate} \\item The reference phase of the FRP, to be obtained for each circle; \\item The opening angle for the FRP, to be obtained as a slowly changing function of time; \\item The focal plane orientation, to be obtained as a slowly changing function of time. \\item The focal plane geometry, to be obtained as a fixed set of parameters. \\end{enumerate} The first three points concern the first order geometric orientation of the field of view: its shifts perpendicular to, and along, the scan direction, and its rotation. These result in systematic shifts in the abscissae and blurred intensity distributions for point sources, both as functions of the field of view position of the detector. The systematic shifts can be solved for as instrument parameters in the positional sphere reconstruction. This also applies to the calibration of the along-scan position of each detector in the focal plane geometry. The opening angle for each detector can only be derived from the brightness distributions of the point sources observed with it. A maximum likelihood solution, optimizing the intensity distributions of the brightest, non-variable point sources should provide these calibration values. This can, however, only be applied to sources for which the positions on the sky are well determined in both coordinates through the positional sphere reconstruction. For the {\\it Planck} mission this condition limits its application to point sources near to the ecliptic poles. Additional information can be obtained from transits of planets and minor planets, for which the absolute coordinates at any time during the mission are known to a much higher accuracy than required for the {\\it Planck} calibrations, and which have the advantage of being visible to most or all detectors. To use these transits for opening angle calibrations requires, however, accurate knowledge of the beam profile perpendicular to the scan direction as well as accurate predictions of their expected intensities, both of which may be difficult to obtain. \\section[]{Iterations with calibrations and the ring analysis} \\label{Sec_iterate} Iterations between the point-source catalogue and the ring reductions are necessary to obtain a properly calibrated geometric reference system for the observations. Without these calibrations in place interpretation of the data in the form of (partial) maps will be of limited value, especially towards the higher frequencies in the power spectrum. Iteration with the BPSC is also required to assign ordinates to point-source transits, which is essential in calibrating the two-dimensional beam profile. Using the BPSC for the identification of point sources in the final ring reductions ensures that the background signals for all all rings contain compatible information. Inconsistent point-source subtraction would lead to harmonic signals in the ring analysis that can not be combined in the harmonic map-making. An iteration with the harmonic map-making \\cite{C1} for the half-year data is required to remove the spectral shift in the data which results from the velocity vector of the satellite (see Section~\\ref{sec_aberr}). While the CMB dipole affects only the CMB component in the frequency maps, the satellite's velocity vector affects all component on every frequency map in both aberration and Doppler shift. Most of these iterations require complete reprocessing of the ring data and are, as such, very time consuming. Without them, however, the scientific interpretation of the data will be subject to considerable uncertainty. ", + "conclusions": "\\label{Sec_concl} The harmonic data model, of which the first part was presented in the current paper, provides a high level of information preservation in the data reductions. It defines calibration requirements and methods and provides a clear path from observed quantities to the scientific products. The latter is important, as the interpretation of the science products requires reliable knowledge of their statistical characteristics, which are well defined in the harmonic model. Although the methods presented in the current paper were developed as part of the harmonic data model, most would also be useful when using the more traditional pixelisation methods. This applies, for example, to the iterative cycle of the ring reductions, BPSC construction and the geometric calibrations. Phase binning of the data can also be used with pixel-based methods, as it provides a means for short-term response calibrations, point source recognition and data compression. Further work is in progress in areas of point-source recognition from the ring data, and the cross-identification of point sources as detected on different rings. The part of the {\\it Planck} data processing presented in this paper will be very demanding. An estimated half a million rings will be produced by the HFI per year of observations. The calibration requirements will make it necessary to process each ring at least three times to get all geometric and beam profile calibrations implemented. The processing of such large quantities of data requires careful planning." + }, + "0112/astro-ph0112089_arXiv.txt": { + "abstract": "As an entry for the 2001 Gordon Bell Award in the ``special'' category, we describe our 3-d, hybrid, adaptive mesh refinement (AMR) code {\\em Enzo} designed for high-resolution, multiphysics, cosmological structure formation simulations. Our parallel implementation places no limit on the depth or complexity of the adaptive grid hierarchy, allowing us to achieve unprecedented spatial and temporal dynamic range. We report on a simulation of primordial star formation which develops over 8000 subgrids at 34 levels of refinement to achieve a local refinement of a factor of $10^{12}$ in space and time. This allows us to resolve the properties of the first stars which form in the universe assuming standard physics and a standard cosmological model. Achieving {\\em extreme resolution} requires the use of 128-bit extended precision arithmetic (EPA) to accurately specify the subgrid positions. We describe our EPA AMR implementation on the IBM SP2 Blue Horizon system at the San Diego Supercomputer Center. ", + "introduction": "Cosmic structure is formed by the gravitational amplification of initially small density fluctuations present in the early universe. A fluctuation containing the mass of the Milky Way galaxy will collapse by a factor of $\\sim 10^3$ in size before it comes into dynamical equilibrium. In order to adequately resolve its internal structure, another two orders of magnitude of spatial resolution {\\em per dimension} are needed, at a minimum, for a total spatial dynamic range $SDR=10^5$. Resolving the formation of individual stars in an entire galaxy would require vastly more resolution: $SDR\\sim 10^{20}$---a seemingly unreachable goal. N-body tree codes are widely used in numerical cosmology to achieve high spatial dynamic ranges in general 3-d evolutions, and some parallel implementations have won past Gordon Bell Awards \\cite{warren92,fukushige96,warren97,makino00}. The largest cosmological N-body simulation used just over $10^9$ particles and achieved a $SDR \\sim 10^4$ \\cite{virgo}. The highest dynamic range N-body simulation achieved $SDR =2\\times10^5$ with substantially fewer particles \\cite{springel00}. These calculations simulate only the collisionless cold dark matter (CDM) which dominates the gravitational dynamics of structure formation. We are interested in simulating the formation of cosmic structures including the all--important baryonic gas which forms the visible galaxies. Our original design goal was to achieve $SDR=10^4$, a mark which we have far surpassed. In this paper we describe the Enzo cosmological adaptive mesh refinement (AMR) code we have developed for parallel computers and its application to a simulation of the formation of the first stars in the universe. The time-dependent calculation is carried out in full 3-d on a structured adaptive grid hierarchy which follows the collapsing protogalaxy and subsequent protostellar cloud to near stellar density starting from primordial fluctuations a few million years after the big bang. More than 8000 subgrids at 34 levels of refinement are generated automatically to achieve a final spatial and temporal dynamic range of $10^{12}$ (for comparison, $10^{12}$ is roughly the ratio of the diameter of the earth to the size of a human cell). It is the highest dynamic range 3-d simulation ever carried out in astrophysics. The calculation incorporates all known relevant dynamical, chemical, and thermodynamic processes and is carried out in a proper expanding cosmological background spacetime. Enzo combines an Euler solver for the primordial gas \\cite{bryan95}, an N-body solver for the collisionless dark matter \\cite{bryan98}, a Poisson solver for the gravitational field \\cite{norman98}, and a 12-species stiff reaction flow solver for the primordial gas chemistry \\cite{aazn97}. The latter is needed to determine the nonequilibrium abundance of molecular hydrogen which dominates the radiative cooling of the gas. This system of equations is solved on every level of the grid hierarchy with full two-way coupling, making it one of the most complex AMR simulations ever carried out. Adaptive mesh refinement allows us to resolve locally all the important length and timescales everywhere within the gravitationally collapsing cloud at all times, giving us confidence that our results are accurate. The calculation is effectively an {\\em ab initio} simulation of star formation which connects initial conditions to the final state in full generality. In this paper, we concentrate on the technical and performance aspects of the AMR code, but also discuss some of the exciting astrophysical results we have obtained. As we will see, achieving extreme resolution is more a matter of high performance data structures and extended precision arithmetic than raw gigaflops, although the latter is certainly needed. ", + "conclusions": "We have presented results on the highest-resolution 3-d computational fluid dynamics simulation yet performed, achieving a spatial dynamic range of $10^{12}$ in the region of interest. The simulation breaks new ground in the complexity of the method and the realism of the result. On the computational front, we solve a series of coupled partial and ordinary differential equations with a range of techniques, in parallel, on a hierarchy of 34 levels containing nearly 10,000 grids. From an astrophysical point of view, we have performed a long sought feat: the {\\it ab initio} simulation of one of the first stars in the universe. Put together, we argue that this represents a watershed in scientific computing. \\vspace{1cm} We acknowledge support from NSF grant AST-9803137. Support for GLB was provided by NASA through Hubble Fellowship grant HF-01104.01-98A from the Space Telescope Science Institute, operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS6-26555. We gratefully acknowledge the support of the San Diego Supercomputing Center." + }, + "0112/astro-ph0112040_arXiv.txt": { + "abstract": "Structural parameters (half-light radius $r_e$, mean effective surface brightness $\\langle\\mu\\rangle_e$, and Sersic index $n$, parameterizing the light profile shape) are derived for a sample of galaxies in the rich cluster AC\\,118 at $z=0.31$: so far the largest (N=93) sample of galaxies at intermediate-redshift with structural parameters measured in the near-infrared. The parameters are obtained in two optical wavebands ($V$ and $R$) and in the $K-$band, corresponding approximately to $B$, $V$ and $H$ rest-frame. The distributions of $r_e$ at $z=0.31$ match those for the Coma cluster (i.e. for the local universe) both in the optical and in the NIR. The $K-$band distribution is of particular interest, since the NIR light mimics the mass distribution of galaxies. The similarity of the distributions for the two clusters (AC\\,118 and Coma) proves that the galaxies at the bright end of the luminosity function did not significantly change their sizes since $z\\sim0.3$ to the present epoch. The ratio of the optical to the NIR half-light radius shows a marked trend with the shape of the light profile (Sersic index $n$). In galaxies with $n\\gtrsim4$ (typical bright ellipticals) $r_{e,NIR}\\sim 0.6 r_{e,opt}$, while the average ratio is 0.8 for galaxies with lower $n$ (typical disk systems). Moreover, the NIR Sersic index is systematically larger than in the optical for $n\\lesssim4$. These results, translated into optical and optical-NIR color gradients, imply that the optical color gradients at $z\\sim$0.3 are similar to those of nearby galaxies. The optical-NIR color gradients are in the average larger, ranging from -0.73 mag/dex for $n\\lesssim4$ to -0.35 mag/dex for $n\\gtrsim4$. Models with `pure age' or `pure metallicity' gradients are unable to reconcile our color gradients estimates with observations at $z\\sim 0$, but we argue that the combined effects of age and metallicity might explain consistently the observed data: passive evolution (plus the possible effect of dust absorption) may account for the differences between the optical and NIR structural properties. The lack of any major change in $r_{e,NIR}$ since $z\\sim0.3$ suggests that merging involving bright galaxies did not play a significant role in the last $\\sim$4.4 Gyr ($\\Omega_M$=0.3, $\\Omega_{\\Lambda}=0$, H$_0=50$ km s$^{-1}$Mpc$^{-1}$). The results of the present paper will be applied to the study of the scaling laws in subsequent works. ", + "introduction": "Our knowledge on the evolution of galaxies comes from the study of the properties of families of galaxies at different redshifts. Observable quantities like magnitudes, colors, structural parameters, velocity dispersions and line strengths are combined into well-established relations like, for instance, the color-magnitude relation (e.g. Bower, Lucey, \\& Ellis 1992; Kodama et al. 1998), the Mg-$\\sigma$ relation \\citep{ZIB97}, the fundamental plane (Dressler et al. 1987; Djorgovski \\& Davies 1987; J\\o rgensen, Franx, \\& Kj\\ae rgaard 1996) and its projection in the $r_e$ -- $\\langle\\mu\\rangle_e$ plane, that is the Kormendy relation (Kormendy 1977; Capaccioli, Caon, D'Onofrio 1992; Ziegler et al. 1999). These relations can now be studied over a significant fraction of the age of the universe. Studies of the properties of early-type galaxies in rich clusters lead to a large amount of observational evidence in favor of the passive evolution scenario, in which an early ($z>$2) and short ($\\sim$1 Gyr) burst of star formation is followed by a long period where the aging of stellar populations drives the observed properties of galaxies through intermediate redshifts until the present epoch (Kodama \\& Arimoto 1997; Stanford, Eisenhardt, \\& Dickinson 1998; J\\o rgensen et al. 1999; Kelson et al. 2000a, and references therein). On the other side, the alternative scenario of the hierarchical merging seems also consistent with most of the existing data (e.g. Kauffmann \\& Charlot 1998), indicating either that the present observational evidence is not sufficient to discriminate between the two pictures or that galaxy evolution is driven by both mechanisms. The half-light radius and the mean surface brightness of a galaxy measure the size and the density of the luminous matter. These quantities are related both to the gravitational potential of the galaxy and to the properties of the stellar populations. The relations existing between the structural parameters of galaxies have been recently studied at intermediate redshifts, both in the form of $r_e$ vs. $\\langle\\mu\\rangle_e$ \\citep[and references therein]{ZSB99} and as luminosity-size relation \\citep{SLL96,SKF99}. \\citet{SLL96} found that disks at $z>$0.5 are in the average $\\sim$ 1.5 mag brighter than in local galaxies. This was contrasted by the results of \\citet{SKF99} who found no luminosity evolution in disks up to $z\\sim$0.9. \\citet{ZSB99} found that the Kormendy relation of cluster galaxies at $z=0.4$ and $z=0.55$ is consistent with the passive evolution scenario. The Kormendy relation has also proven to be a powerful tool to test the expansion of the universe by the measurement of the cosmological dimming (i.e. for the Tolman test). \\citet{LUB01} studied the Kormendy relation up to $z=0.92$, finding that the data are consistent with the cosmic expansion together with passive evolution. Finally, the structural parameters enter the fundamental plane of early-type galaxies, of which the Kormendy relation is a two-dimensional projection. Due to its very small scatter, the fundamental plane is a powerful tool to study the evolution of the M/L ratio of galaxies belonging to distant clusters. The recent determinations of the fundamental plane at intermediate redshifts constrain the epoch of star formation in early-type galaxies at large redshifts ($z>$2 for $\\Omega_M$=0.3 or $z>$5 for $\\Omega_M$=1 and $\\Omega_{\\Lambda}=0$; see van Dokkum \\& Franx 1996; J\\o rgensen et al. 1999; Kelson et al. 2000a). However, these studies have to face with the problem of the progenitor bias, i.e. with our ignorance of the relationship between local galaxies and their high redshift counterparts (see van Dokkum \\& Franx 2001 for a thorough discussion). A clue for the interpretation and for the applications of the scaling laws is the dependence of the structural parameters on wavelength. The basic distinction between optical and NIR is that the optical wavebands are very sensitive to the effects of metallicity (through line-blanketing) and to the content of younger populations (dust absorption included), while the NIR wavebands are dominated by the old, more quiescent, stellar populations and are virtually unaffected by dust absorption. For these reasons the NIR is the best tracer of the mass distribution inside a galaxy. Local galaxies are smaller at NIR wavebands than in the visible, with a similar, though weaker, trend also within the optical wavelength range (Scodeggio et al. 1998; Pahre, de Carvalho, \\& Djorgovski 1998). This behavior is confirmed by the relatively large optical-NIR color gradients measured in early-type galaxies, the optical-optical gradients being usually much smaller, except when the $U-$band is involved (Peletier et al. 1990a; Peletier, Valentijn, \\& Jameson 1990b; Hinkley $\\&$ Im 2001; Nelson et al. 2001). Small optical color gradients were recently measured in distant (up to $z\\sim$1) cluster early-type galaxies by \\citet{SMG00} and by \\citet{TKA00}. These authors concluded that metallicity is the primary factor driving the color gradients. The derivation of the structural parameters at intermediate redshifts is affected by the problem of the small angular size of the galaxies, whose half-light radii are typically less than 1 arcsec. For this reason many works were based on imaging taken with the HST. It has been demonstrated, however, that reliable structural parameters can be derived from ground-based photometry as well, on high resolution ($< 0.1 ''$/pxl) and exceptional seeing ($<0.4''$) conditions, provided that the properties of the PSF are carefully taken into account \\citep{SLL96,JOR99}. In this paper we obtain the first multi-band large set of structural parameters for galaxies at an intermediate redshift ($z=0.31$). The data are derived from ground-based (ESO-NTT) photometry in two optical ($V$ and $R$) and one NIR ($K$) wavebands. The derived values for the structural parameters will be published in a forthcoming paper addressed to the study of galaxy evolution by means of the scaling laws (Busarello et al. 2001a, in preparation). The paper is organized as follows. The sample and the data used in this study are presented in \\S~\\ref{samples}. The derivation of the structural parameters is outlined in \\S~\\ref{STRSEC}. The structural parameters from HST imaging are derived in \\S~\\ref{HST_HST}, where the results from Sersic models and bulge-disk decomposition are compared. In \\S~\\ref{SEC_HST_R} we use HST photometry for a subset of galaxies to verify the reliability of ground-based structural parameters. The structural parameters are presented in \\S~\\ref{OP_NIR}, where we make a comparison between optical and NIR properties at $z\\sim0.3$ and with the values at $z\\sim0$. In \\S~\\ref{SECGRAD} we analyze the optical and NIR structure of galaxies in terms of internal color gradients. The results are discussed in \\S~\\ref{CONC}. ", + "conclusions": "\\label{CONC} In this study we derived optical and NIR structural parameters for galaxies belonging to the cluster AC\\,118 at $z=0.31$ from ground-based photometry. The present data set constitutes the first large sample of structural parameters determined at intermediate redshifts in different wavebands and, in particular, it is the first such sample in the near-infrared. The samples consist in $N$=93 objects studied in the $R$ and $K$ wavebands and in $N$=58 galaxies studied in the $V-$band. The reliability of the data was verified by comparison with the results from HST photometry for a subsample of galaxies: the scatter between measurements from HST- and from ground-based photometry is equivalent to the scatter between parameters determined with two different fitting methods from HST images, namely of $\\sim$35\\% in $r_e$. Given that these values of the scatter are generally claimed to be a measure of the `intrinsic uncertainty' in structural parameter measurements, we can state that reliable structural parameters can be still derived from ground-based photometry in ordinary observing conditions at least up to $z\\sim0.3$. Since morphology was not a selection criterion in this study, some remarks are needed for the comparison of our results with other works. \\citet{CBS98} established that $78\\%$ of the galaxies in the core of AC\\,118 are early-types while $22\\%$ are spirals. They also found a significant rise in the fraction of S0 galaxies with respect to the local density-morphology relation. Morphological classification from \\citet{CBS98} is available for 21 galaxies in sample D. The number of spirals is 5, of which 3 have $n<$1.5 while the other two have $n \\gtrsim 4$. Only one of the spirals has `late-type' colors, while the others lie on the red sequence of AC\\,118. It is interesting to notice that all the remaining galaxies with `late-type' colors lie outside of the cluster core and have $n < 2$. They are likely disk galaxies dominated by young stellar populations. These facts are in agreement with a scenario in which spirals are accreted into the cluster from the field, cease to form stars, become gas-deficient objects and eventually undergo a morphological transformation \\citep{CBS98}. The remaining galaxies with morphological information consist in 10 S0, eight of which with $2\\lesssim n \\lesssim 3$ and two with $n \\sim 8$, and 6 ellipticals, for which we find $n \\gtrsim 4$. We stress that the galaxies with large disks and late-type colors are only $4\\%$ of galaxies in sample D. The main feature concerning the multi-band properties of galaxies is that these are much more concentrated in the NIR than in the optical. The mean ratio of NIR to optical half-light radii ($\\sim 0.75$) is not different from that found in local samples of early-type galaxies ($\\sim 0.8$, e.g. Pahre 1999), but the situation is complicated by the systematic differences found in the shape of the brightness distribution: galaxies with low $n$ value have a light profile more peaked in the center in the $K$-band ($n_K - n_R \\sim 1$), while for $n \\gtrsim 4$ a lower $r_{e,K}/r_{e,R}$ is observed ($\\sim 0.6$). Such trends can be interpreted as due to the increase of the disk fraction when $n$ approaches unity. To discuss the combined effects of half-light radii and Sersic indices, we derived the optical and NIR color gradients by means of equation 4. Most of galaxies in AC\\,118 do not show significant optical color gradients, with a median value of $-0.06^{+0.03}_{-0.15}$. This result is consistent with previous estimates of optical color gradients at intermediate redshifts \\citep{SMG00,TAO00}. Conversely, strong negative gradients are found in optical--NIR colors: the median of the $R\\!-\\!K$ gradients amounts to $-0.49^{-0.43}_{-0.55}$. Moreover, this effect depends on the shape of the light profile: the $R\\!-\\!K$ gradient increases from $-0.73^{-0.63}_{-0.83}$ for galaxies with low $n$ value to $-0.35^{-0.29}_{-0.42}$ for $n \\gtrsim 4$. Again this result can be explained by an increase in the disk fraction when $n \\rightarrow 1$. To interpret these results in terms of the properties of the stellar populations, we introduced a pure age and a pure metallicity model. We used the $R\\!-\\!K$ gradients to calibrate both models and to predict $V\\!-\\!R$ color gradients at $z=0.31$. Moreover, the evolution of color gradients as inferred by the models was compared with the measurements by \\citet{PDD90,PVJ90} for a sample of nearby elliptical galaxies. The age model is not able to reconcile our optical and NIR estimates at $z=0.31$ since a significant negative gradient is predicted in $V\\!-\\!R$ ($\\sim -0.26$). On the other hand, the predicted color gradient agree with the measurements at $z=0$. Conversely, the pure metallicity model gives correct results at $z=0.31$ but predicts too strong NIR gradients at $z=0$. Since the age model is able to describe the evolution of color gradients, whereas the metallicity model correctly explains the dependence of color gradients on wavelength, we argue that a model in which effects of age and metallicity are consistently combined could explain optical and NIR color gradients from $z=0.31$ up to $z=0$. Finally, we remark that dust likely affects the optical-NIR color gradients. Dust is actually presents in the core of most early-type galaxies \\citep{GDJ95,VDF96} and is known to produce high central color gradients in bulges of S0 and spiral galaxies (e.g. Peletier et al. 1999). The distributions of both the optical and NIR half-light radii of galaxies at $z\\sim0.3$ are found to be consistent with those of Coma early-type galaxies. The result for the NIR waveband is the most significant, since the NIR probes directly the mass distribution inside galaxies. The observed distributions thus indicate that the sizes of bright galaxies did not substantially change since $z\\sim 0.3$ to the present epoch. This result suggests that (dissipation-less) merging did not play a major role in galaxy evolution during the last $4.4$ Gyr (see also Nelson et al. 2001)." + }, + "0112/astro-ph0112262_arXiv.txt": { + "abstract": "We present the results of detailed studies of the astrophysical conditions in $z\\sim 3$ Lyman Break Galaxies (LBGs), placing particular emphasis on what is learned from LBG rest frame UV spectra. By drawing from our database of $\\sim 1000$ spectra, and constructing higher S/N composite spectra from galaxies grouped according to properties such as luminosity, extinction, morphology, and environment, we can show how the rest-frame UV spectroscopic properties systematically depend on other galaxy properties. Such information is crucial to understanding the detailed nature of LBGs, and their impact on the surrounding IGM. ", + "introduction": "Until now, Lyman Break Galaxy (LBG) rest-frame UV spectra have been primarily used to measure redshifts. The measured redshifts confirm the high-redshift nature of galaxy candidates selected by their distinctive broadband optical colors (Steidel et al. 1996); enable the study of the spatial clustering of these galaxies (Adelberger et al. 1998); and, when combined with the apparent magnitudes and colors of LBGs, can be used to construct the rest-frame UV luminosity function and UV luminosity density at $z\\sim 3$ (Steidel et al. 1999). ", + "conclusions": "Rapid star-formation has a profound effect on the ISM of LBGs, as seen in the large absorption equivalent widths and significant velocity offsets between rest-frame UV emission and absorption features. The effects of star-formation also extend to the surrounding inter-galactic medium, through the leakage of Lyman continuum emission, and from shock heating by outflowing material. A systematic analysis of the rest-frame UV spectroscopic of LBGs will help us understand how star-formation transforms both galaxies and their surrounding environment at $z\\sim 3$." + }, + "0112/astro-ph0112438_arXiv.txt": { + "abstract": "Quintessence, a time-varying energy component that may account for the accelerated expansion of the universe, can be characterized by its equation of state and sound speed. In this paper, we show that if the quintessence density is at least one percent of the critical density at the surface of last scattering the cosmic microwave background anisotropy can distinguish between models whose sound speed is near the speed of light versus near zero, which could be useful in distinguishing competing candidates for dark energy. ", + "introduction": " ", + "conclusions": "" + }, + "0112/astro-ph0112112_arXiv.txt": { + "abstract": "We present results from monitoring of the distant (z = 2.64), gravitationally lensed quasar MG~J0414+0534 with the {\\it Chandra X-ray Observatory}. An Fe~K$\\alpha$ line at 6.49 $\\pm$ 0.09~keV (rest-frame) with an equivalent width of $\\sim$ 190~eV consistent with fluorescence from a cold medium is detected at the 99\\% confidence level in the spectrum of the brightest image A. During the last two observations of our monitoring program we detected a five-fold increase of the equivalent width of a narrow Fe~K$\\alpha$ line in the spectrum of image B but not in the brighter image A whereas image C is too faint to resolve the line. The continuum emission component of image B did not follow the sudden enhancement of the iron line in the last two observations. We propose that the sudden increase in the iron line strength from $\\sim$ 190~eV to 900~eV only in image B can be explained with a caustic crossing due to microlensing that selectively enhances a strip of the line emission region of the accretion disk. The non-enhancement of the continuum emission in the spectrum of image B suggests that the continuum emission region is concentrated closer to the center of the accretion disk than the iron line emission region and the magnification caustic has not reached close enough to the former region to amplify it. A model of a caustic crossing event predicts discontinuities in the light-curve of the magnification and provides an upper limit of $\\sim$ 5 $\\times$ 10$^{-4}$ pc on the outer radius of the Fe~K$\\alpha$ emission region. The non-detection of any relativistic or Doppler shifts of the iron line in the spectrum of image B implies that the magnification caustic for the last two observations was located at a radius greater than $\\sim$ 100 gravitational radii. Each observation of the quadruply lensed quasar MG~J0414+0534 provides a view of the quasar at four different epochs spaced by the time-delays between the lensed images. We produced a light-curve of the quasar X-ray flux by normalizing the flux of each image to the mean flux of that image over all observations. We find significant deviations of the normalized light-curve from unity especially for the faintest image C. A plausible mechanism to explain the flux variability of image C is a microlensing event. Finally, spectral analysis of MG~J0414+0534 indicates the presence of significant absorption in excess of the Galactic value. For absorption at the redshift of the lensed quasar we find an intrinsic column density of N$_{H}$ $\\sim$ 5 $\\times$ 10$^{22}$ cm$^{-2}$, consistent with the reddening observed in the optical band. ", + "introduction": "Direct imaging of the immediate environments of black holes in Active Galactic Nuclei (AGNs) is beyond the capabilities of present day telescopes; rough estimates of the characteristic sizes of emission regions of AGN in the radio, optical and X-ray bands, based on light-travel time arguments, correspond to angular sizes on the order of tens of nano-arcseconds at a redshift of $z=1$. However, indirect mapping methods have been developed recently to study the environs of AGN without the need of nano-arcsec resolution. One of these methods, known as reverberation mapping (Blandford \\& McKee 1982), has been successfully used to determine the size and probe the structure of the broad-line region in several AGN (see Peterson 1993 and Netzer \\& Peterson 1997 for recent reviews). This method relies on the time lag between variations in the flux from a central source of ionizing radiation and the response of the emission lines from photo-ionized gas in the broad-line region. A closely related technique for indirect mapping of the accretion disk around the black hole relies on the variation of the profiles of the Fe~K$\\alpha$ fluorescence lines originating in the inner parts of the disk. The line profiles are determined by a combination of the Doppler effect and special and general relativistic effects (e.g., Fabian et al 1989; Laor 1991) while their variations are caused by fluctuations in the X-ray continuum that drive the line emission (Young \\& Reynolds 2000). A recently proposed, independent, method for indirect imaging of AGN accretion disks exploits the high-magnification microlensing events (HMEs) in gravitational lens (GL) quasars (Grieger et al. 1988 and 1991; Schneider, Ehlers \\& Falco 1992; Gould \\& Gaudi 1997; Agol \\& Krolik 1999; Yonehara et al. 1999; Mineshige \\& Yonehara 1999, Yonehara 2001; Shalyapin 2001). During a microlensing event magnification caustics produced by stars in the lensing galaxy traverse the plane of the accretion disk and selectively magnify different emission regions. An analysis of the light-curves of microlensing events obtained in several wavelengths can be used to infer the surface brightness and inclination angle of the accretion disk and possibly provide constraints on the mass and spin of the black hole (e.g, Agol \\& Krolik 1999 and references within). A promising object for application of this technique is the Einstein Cross. With the first detection of a microlensing event in the Einstein Cross by Irwin et al. (1989) and subsequent verification by Corrigan et al. (1991) and Ostensen et al. (1996) it became clear that future monitoring of the Einstein Cross has the potential of yielding information on the geometry of the accretion disk and possibly the spin and mass of the black hole. Thus, a program to monitor magnification events in the Einstein Cross has been undertaken as part of the Optical Gravitational Lensing Experiment (OGLE, see Wo{\\' z}niak et al. 1999 and references within). Microlensing events have been detected with OGLE, however a clear case of a HME due to a caustic crossing has yet to be detected. Recently, constraints on the continuum source size of the lensed quasar have been derived based on the microlensing light-curves of the Einstein Cross (eg., Wyithe, Webster \\& Turner 2000; Yonehara 2001; Shalyapin 2001). In this paper we report the serendipitous detection of significant variations in the shape and equivalent width (EW) of a reprocessed Fe~K${\\alpha}$ line only in one of the lensed images of the quasar MG~J0414+0534. We interpret this change as the result of the crossing of the accretion disk by a magnification caustic. The gravitational lens system MG~J0414+0534 was discovered by Hewitt et al. (1992). The system consists of four lensed images commonly referred to as A1, A2, B and C (see Figure 1) with the separation between A1 and A2 being $\\sim$ 0{\\sarc}4. The lensed source is a distant (z = 2.639), low-luminosity radio quasar (Lawrence et al. 1995). A faint blue arc is detected in HST observations (Falco et al. 1997) that extends between images A1 and B. The arc is most likely the host galaxy of the quasar lensed by the foreground elliptical galaxy located at a redshift of 0.96 (Tonry \\& Kochanek 1999). An object located $\\sim$ 1~{\\arcsec} West of image B and referred to as object X (Schechter \\& Moore 1993) possibly contributes to the lensing effect. IR observations of MG~J0414+0532 indicate that the components of this system are exceedingly red with respective R-H colors of 6.8 mag for image C, 3.2~mag for the arc, and 3.2 mag for the lens galaxy. The origin of the absorption is controversial with early studies mostly favoring the dusty lens hypothesis (Lawrence et al. 1995; Malhotra, Rhoad, \\& Turner 1997; McLeod et al. 1998) and more recent work suggesting a host galaxy origin (Tonry \\& Kochanek 1999; Falco et al. 1997) or a combination of lens and host (Angonin-Willaime et al. 1999). The flux ratios have been observed to be a function of wavelength and this has been interpreted as either due to extinction in the lens or microlensing of one or more of the images. Recent VLBI and HST NICMOS H-band images have been used to produce accurate lens models for this system (Ros et al. 2000). These models predict the time delays between the lensed images to be ${\\Delta}t_{AB}$ $\\sim$ 16 $\\pm$ 1.4 days and ${\\Delta}t_{CB}$ $\\sim$ 66 $\\pm$ 5 days. The time-delays predicted from the lens model of Ros et al. (2000) assume a flat cosmology with $\\Omega_{0}=1$ and $H_{0} = 65$ km s$^{-1}$ Mpc$^{-1}$. The monitoring program of MG~J0414+0532 with {\\it Chandra} was planned such that the time interval between observations was substantially shorter than a year, roughly the predicted average rise for a high magnification event due to microlensing in this GL system (Witt, Mao \\& Schechter 1995). A detailed description of the {\\it Chandra} observations of MG~J0414+0534 and the data reduction procedures used is presented in section 2. The resolved X-ray lensed images of this GL system are presented in section 3. The spectral analysis revealing the properties of an intervening absorber and reprocessed iron line features is detailed in section 4. We conclude with a discussion of the origin of the detected iron line features and the prospects of mapping the accretion disk of a distant quasar with microlensing of a reprocessed Fe~K${\\alpha}$ line. ", + "conclusions": "" + }, + "0112/astro-ph0112442_arXiv.txt": { + "abstract": "{Narrowband spikes of the decimeter type have been identified in dynamic spectrograms of Phoenix-2 of ETH Zurich and located in position with the Nan\\c{c}ay Radioheliograph at the same frequency. The spike positions have been compared with the location of hard X-ray emission and the thermal flare plasma in soft X-rays and EUV lines. The decimetric spikes are found to be single sources located some 20\" to 400\" away from the flare site in hard or soft X-rays. In most cases there is no bright footpoint nearby. In at least two cases the spikes are near loop tops. These observations do not confirm the widely held view that the spike emission is produced by some loss-cone instability masering near the footpoints of flare loops. On the other hand, the large distance to the flare sites and the fact that these spikes are all observed in the flare decay phase make the analyzed spike sources questionable sites for the main flare electron acceleration. They possibly indicate coronal post-flare acceleration sites. ", + "introduction": "The acceleration of large numbers of particles in flares is an old, major enigma in solar physics. Not only the question concerning the predominant mechanism is unsolved, even the location of acceleration is unclear. Nevertheless, over the past decade crucial observations on small relatives of flares have accumulated. They evidence that short narrowband radio bursts in meter wavelengths are signatures of the acceleration process. Benz et al. (1996) have shown that narrowband {\\sl metric spikes} in general correlate with type III bursts starting at slightly lower frequencies. The metric spikes have been found by Paesold et al. (2001) to be located generally on the extension of type III trajectories and suggest a model of energy release taking place in or close to the spike sources. Modeling coronal densities (Paesold et al. 2001) and spatially resolved observations of metric spike events (Krucker et al. 1995; 1997) put the sources at altitudes of $2\\times 10^{10}$cm and more. Similarly, {\\sl type I} radio bursts in noise storms have been found e.g. by Raulin et al. (1994) to be accompanied by a delayed continuum radio emission and a soft X-ray brightening. The similarity of spikes and type I bursts has been emphasized previously (Benz 1985). Both noise storms and metric spike/type III bursts are not associated with regular, hard X-ray and centimeter-wave emitting flares. The short duration and the narrow bandwidth suggest a small source size and therefore a high radio brightness temperature (order of $10^{15}$ K). Only a coherent mechanism can account for the emission, but none of the published mechanisms is generally accepted. A model proposed by Benz \\& Wentzel (1981) suggests type I radiation to originate in the acceleration region, where waves driven by unstable currents couple with Langmuir waves produced by accelerated electron beams. Similar processes have been proposed for narrowband spikes in the decimeter range (cf. below). The situation is much less clear for narrowband spikes at {\\sl decimeter} wavelength. The term 'narrowband, millisecond spikes' has been introduced for them referring to narrowband (few percent of the center frequency) and short (few tens of ms) peaks above 1 GHz (Droege 1977; Slottje 1978). Originally, they were reported to occur in the rise phase of centimeter radio bursts and thus to be associated with major flares. The association rate with hard X-ray flares is high (Benz \\& Kane 1986; G\\\"udel et al. 1991). However, a delay of spike groups and hard X-ray peaks of the order of a few seconds has been noted (Aschwanden \\& G\\\"udel 1992). Moreover, spikes have been discovered also in decimeter type IV bursts occurring after the HXR emitting phase of flares (Isliker \\& Benz 1994). Contrary to their relatives at meter waves, decimetric spikes do not correlate with type III bursts. The emission process of decimetric spikes is highly controversial. Originally, a loss-cone instability of trapped electrons has been proposed to produce electron cyclotron maser emission at the footpoints of flare loops (Holman et al. 1980; Melrose \\& Dulk 1982). To avoid the assumption of high magnetic field strength in the source, the model has been changed to emission of upper-hybrid and bernstein modes (Willes \\& Robinson 1996). The scheme can interpret occasional harmonic emission in decimeter spikes (Benz \\& G\\\"udel 1987; Krucker \\& Benz 1994). Alternatively, Tajima et al. (1990) and G\\\"udel \\& Wentzel (1993) proposed the spike sources to be in the acceleration regions of flares and to result from waves produced by the acceleration process. Here we set out to test the predictions of the two emission models concerning the source position of narrowband decimeter spikes. In the loss-cone scenario, the spike emission is a secondary phenomenon located near the footpoints of flaring loops. Hard X-ray sources, produced by precipitating electrons have been found at footpoints of such loops (Duijveman et al. 1982). Thus the loss-cone scenario predicts {\\sl (i)} spikes to occur at similar locations when seen from Earth. {\\sl (ii)} As hard X-rays often have double sources, the spikes may behave similarly (as proposed by Conway \\& Willes 2000). In the acceleration-region scenario on the other hand, the spikes are expected to originate from single sources at some distance from the footpoints, at higher altitude, possibly at the top of loops, and to be single sources. Previously, Gary et al. (1991) have reported imaging observations of possible decimetric spikes at 2.8 GHz originating some 25\" from the gyrosynchrotron source. An observation at 5.7 GHz of a limb flare reported by Altyntsev et al. (1995) suggests the spikes to be located within the region of gyrosynchrotron emission about 35 Mm above the photosphere. These observations did not have spectral resolution to confirm the narrowbandedness of the emission. As there are several other types of short duration coherent emission in the decimeter range, the identification of these observations as narrowband spikes remains controversial. Here we start out from spike events that are well identified by spectrometer observations and locate their position with simultaneous measurements by an interferometer observing at the frequencies of spike emission. The position of the radio emission is then compared with imaging observations in hard X-rays, if available. Finally, the sources are put into the context of the thermal coronal plasma and field geometry as imaged by coronal EUV lines and soft X-rays. ", + "conclusions": "The requirement to choose decimetric spike events at low frequency to be observable with the NRH has produced a selection of five events that all occurred in the flare decay phase. The analysis indicates that 1. Generally, spike sources are not found at footpoints of flare loops as indicated by soft X-ray or coronal EUV emissions. Only one case (Fig.7) seems to be nearer to footpoints than loop tops, but the footpoints are not those of the main flare loops. There are two clear cases (Figs. 4 and 9) and three possible cases (two in Fig. 3 and one in Fig. 5) of a coincidence with a loop top in projection. 2. Consistent with this, we find the spike sources at very distant locations from the site of hard X-ray emission (Fig.4). The angular distance from the main soft X-ray or EUV loops is 20\" to 400\". Only in one case (Fig.5), does the spike sources coincide with strong soft X-ray emission, which may be a projection effect. 3. None of the spike sources was found to be double at a given time. 4. The spike sources, if measured at different frequencies, appear to be located at the same site within the intrinsic scatter they have in time (Figs. 4, 7, and 9). There is however a clear exception (Fig.5), where the positions are spread according to frequency. 5. There may be more than one source of spike emission in the course of an event. In Fig. 3 a case is shown where two clusters separated by 16 minutes in time originate from widely different places in the same active region. 6. Other radio emissions at the same frequency may originate at far away positions, such as observed for simultaneous type III bursts (Fig. 6) and decimetric pulsations (Fig. 8). The fact that the spike sources are found far away from the major flare site may be related to the low frequency events we have selected. Combined with their occurrence after the main flare phase, they are unlikely to be the main sites of flare electron acceleration. Nevertheless, these spike sources cannot be excluded as secondary acceleration sites as they have all been associated with 2.8 GHz radiation of potential gyrosynchrotron origin. Thus we suggest that the observed spikes are related to the rearrangement of the coronal magnetic field after the main flare at lower altitude. They may be indicative of post-flare high-corona acceleration sites. Narrowband spikes at decimetric frequencies have been reported in the flare rise phase and main phase, even at low frequencies (Benz 1985). It will be interesting to investigate the above findings on the location of spikes earlier in the flare and at higher decimetric frequencies. More sensitive and precise information on hard X-rays are urgently needed. Such a work is well suited for the pending HESSI mission." + }, + "0112/astro-ph0112168_arXiv.txt": { + "abstract": "We discuss the possible connection between supernova explosions (SN) and gamma-ray bursters (GRB) from the perspective of our current understanding of SN physics. Recent evidence strongly suggests that the explosion mechanism of core collapse SN is intrinsically aspherical. Typically, a neutron star is formed. However, the observed properties of the expanding SN envelopes remnants make these objects very unlikely candidates for GRBs. Most candidates for a GRB/SN connection seem to require the prompt or delayed formation of a black hole. These include the collapse of very massive stars (e.g. hypernovae) and 'classical' SNe with a significant fallback of material over time scales of hours to days, resulting in the collapse of the neutron star to a black hole. We suggest the merger of a neutron star with a white dwarf as a subclass of thermonuclear SNe and a potential candidate for a SN/GRB connection. ", + "introduction": "Core collapse supernovae are the final stages of stellar evolution in massive stars during which the central region collapses, forms a neutron star (NS) or a black hole, and the outer layers are ejected. Recent explosion scenarios assumed that the ejection is due to energy deposition by neutrinos into the envelope but detailed models do not produce powerful explosions and, in most cases, do not trigger the ejection of the envelope at all (e.g. Yamada et al. 1999, Ramp \\& Janka 2000, Mezzacappa et al. 2001). There is new and mounting evidence for an asphericity and, in particular, for axial symmetry in several SNe which may be hard to reconcile within the spherical picture. This evidence includes the observed high polarization and its variation with time (e.g. Wang et al. 2001, Leonard et al. 2000), pulsar kicks (Strom et al. 1995), high velocity iron-group and intermediate-mass element material observed in remnants (Lucy 1987, Tueller et al. 1991), direct observations of the debris of SN1987A (Wang et al. 2001, H\\\"oflich, Khokhlov \\& Wang 2001), etc. To be in agreement with the observations, any successful mechanism must invoke some sort of axial symmetry for the explosion. As a limiting case, we consider jet-induced/dominated explosions of \"classical\" core collapse SNe. The discovery of magnetars (Kouveliotou et al. 1998, Duncan \\& Thomson 1992) revived the idea that a MHD-jet with appropriate properties may be formed at the NS or BH (LeBlanc \\& Wilson 1970, Symbalisty 1984). Our study is based on detailed 3-D hydrodynamical and radiation transport models (Khokhlov 1998, H\\\"oflich et al. 1998, 2001). We demonstrate the influence of the jet properties and of progenitor structure on the final density, chemical structure and the fallback of material. \\subsection{Results for Jet-Induced Supernovae} \\noindent {\\bf The Setup:} The computational domain is a cube of size $L$ with a spherical star of radius $R_{\\rm star}$ and mass $M_{\\rm star}$ placed in the center. The innermost part with mass $M_{\\rm core} \\simeq 1.6 M_{\\odot}$ and radius $R_{\\rm core} = 4.5 \\times 10^8$~cm, consisting of Fe and Si, is assumed to have collapsed on a timescale much faster than the outer, lower-density material. It is removed and replaced by a point gravitational source with mass $ M_{\\rm core}$ representing the newly formed neutron star. The remaining mass of the envelope $M_{env}$ is mapped onto the computational domain. At two polar locations where the jets are initiated at $R_{core}$, we impose an inflow with velocity $v_j$ and a density $\\rho_j$. At $R_{\\rm core}$, the jet density and pressure are the same as those of the background material. For the first 0.5~s, the jet velocity at $R_{\\rm core}$ is kept constant at $v_j$. After 0.5~s, the velocity of the jets at $R_{\\rm core}$ was gradually decreased to zero at approximately 2~s. The total energy of the jets is $E_j$. These parameters are consistent with, but somewhat less than, those of the LeBlanc-Wilson model. \\noindent {\\bf The reference model:} As a baseline case, we consider a jet-induced explosion in a helium star. Jet propagation inside the star is shown in Fig. \\ref{jet1}. As the jets move outwards, they remain collimated and do not develop much internal structure. A bow shock forms at the head of the jet and spreads in all directions, roughly cylindrically around each jet. The jet-engine has been switched off after about 2.5 seconds, the material of the bow shock continues to propagate through the star. The stellar material is shocked by the bow shock. Mach shocks travels towards the equator resulting in a redistribution of the energy. The opening angle of the jet depends on the ratio between the velocity of the bow shock to the speed of sound. For a given star, this angle determines the efficiency of the deposition of the jet energy into the stellar envelope. Here, the efficiency of the energy deposition is about 40 \\%, and the final asymmetry of the envelope is about two. \\noindent {\\bf Influence of the jet properties:} Fig. \\ref{jet2} shows two examples of an explosion with a low and a very high jet velocity compared to the baseline case (Fig. \\ref{jet1}). Fig. \\ref{jet2} demonstrates the influence of the jet velocity on the opening angle of the jet and, consequently, on the efficiency of the energy deposition. For the low velocity jet, the jet engine is switched off long before the jet penetrates the stellar envelope. Almost all of the energy of the jet goes into the stellar explosion. On a contrary, the fast jet (61,000 km/sec) triggers only a weak explosion of 0.9 foe although its total energy was $\\approx 10 foe $. \\noindent {\\bf Influence of the progenitor:} For a very extended star, as in case of 'normal' Type II Supernovae, the bow shock of a low velocity jet stalls within the envelope, and the entire jet energy is used to trigger the ejection of the stellar envelope. In our example (Fig. \\ref{model}), the jet material penetrates the helium core at about 100 seconds. After about 250 seconds the material of the jet stalls within the hydrogen rich envelope and after passing about 5 solar masses in the radial mass scale of the spherical progenitor. At this time, the isobars are almost spherical, and an almost spherical shock front travels outwards. Consequently, strong asphericities are limited to the inner regions. After about 385 seconds, we stopped the 3-D run and remaped the outer layers into 1-D structure, and followed the further evolution in 1-D. After about 1.8 $10^ 4$ seconds, the shock front reaches the surface. After about 3 days, the envelope expands homologously. The region where the jet material stalled, expands at velocities of about 4500 km/sec. \\noindent {\\bf Fallback:} Jet-induced SN have very different characteristics with respect to fallback of material and the innermost structure. In 1-D calculations and for stars with main sequence masses of less than 20 $M_\\odot$ and explosion energies in excess of 1 foe, the fallback of material remains less than 1.E-2 to 1.E-3 $M_\\odot$ and an inner, low density cavity is formed with an outer edge of $^{56}Ni$. For explosion energies between 1 and 2 foe, the outer edge of the cavity expands typically with velocities of about 700 to 1500 km/sec (e.g. Woosley 1997, H\\\"oflich et al. 2000). In contrast, we find strong, continuous fallback of $\\approx 0.2 M_\\odot$ in our 3-D models, and no lower limit for the velocity of the expanding material (Fig. 4 of Khokhlov \\& H\\\"oflich 2001). This significant amount of fallback must have important consequences for the secondary formation of a black hole. The exact amount and time scales for the final accretion on the NS will depend sensitively on the rotation and momentum transport. \\noindent {\\bf Chemical Structure:} The final chemical profiles of elements formed during the stellar evolution such as He, C, O and Si are 'butterfly-shaped' whereas the jet material fills an inner, conic structure (Fig. \\ref{model}, upper, middle panel). The composition of the jets must reflect the composition of the innermost parts of the star, and should contain heavy and intermediate-mass elements, freshly synthesized material such as $^{56}Ni$ and, maybe, r-process elements because, in our examples, the entropy at the bow shock region of the jet was as high as a few hundred, a crucial condition for a successful r-process in a high-$Y_e$ environment. In any case, during the explosion, the jets bring heavy and intermediate mass elements into the outer H-rich layers. \\noindent {\\bf Radiation Transport Effects:} For the compact progenitors of SNe~Ib/c, the final departures of the iso-density contours from sphericity are typically a factor of two. This will produce a linear polarization of about 2 to 3 \\% (H\\\"oflich et al. 1995) consistent with the values observed for SNe~Ib/c. In case of a red supergiant, i.e. SNe~II, the asphericity is restricted to the inner layers of the H-rich envelope. There the iso-densities show an axis ratio of up to $\\approx$ 1.3. The intermediate and outer H-rich layers remain spherical. Thus, even within the picture of jet-induced explosion, the latter effect alone cannot (!) account for the high polarization produced in the intermediate H-rich layers of core-collapse SN with a massive envelope such as SN~1999em where asymmetric ionization by $\\gamma $ photons for an aspherical chemical distribution of $^{56}Ni$ is crucial (see Fig. \\ref{model} and H\\\"oflich et al. 2001). \\noindent", + "conclusions": "" + }, + "0112/astro-ph0112397_arXiv.txt": { + "abstract": "We present multi-instrument optical observations of the High Energy Transient Explorer (HETE-2)/Interplanetary Network (IPN) error box of GRB~010921. This event was the first gamma ray burst (GRB) localized by HETE-2 which has resulted in the detection of an optical afterglow. In this paper we report the earliest known observations of the GRB010921 field, taken with the 0.11-m Livermore Optical Transient Imaging System (LOTIS) telescope, and the earliest known detection of the GRB010921 optical afterglow, using the 0.5-m Sloan Digital Sky Survey Photometric Telescope (SDSS PT). Observations with the LOTIS telescope began during a routine sky patrol 52 minutes after the burst. Observations were made with the SDSS PT, the 0.6-m Super-LOTIS telescope, and the 1.34-m Tautenburg Schmidt telescope at 21.3, 21.8, and 37.5 hours after the GRB, respectively. In addition, the host galaxy was observed with the USNOFS 1.0-m telescope 56 days after the burst. We find that at later times (t > 1 day after the burst), the optical afterglow exhibited a power-law decline with a slope of $\\alpha = 1.75 \\pm 0.28$. However, our earliest observations show that this power-law decline can not have extended to early times (t < 0.035 day). ", + "introduction": "The High Energy Transient Explorer (space.mit.edu/HETE/) HETE-2 is dedicated to the study of gamma-ray bursts (GRBs). HETE-2 is currently the only GRB detector capable of localizing and disseminating GRB coordinates in near real-time. Low-energy emission during and shortly after a GRB ($t \\lesssim 1$~hr) potentially holds the key to significant progress in understanding the central engine of GRBs and could provide valuable clues to their progenitors \\citep{meszaros01}. The HETE-2 detection of GRB~010921 together with data from the Interplanetary Network (IPN) provided the first HETE-2 localization which has resulted in the detection of an optical afterglow. Although the afterglow was relatively bright, early observations of the error box failed to reveal any candidate afterglows because of source confusion with its bright host galaxy ($R \\sim 21.7$) \\citep{price01a}. Spectroscopy of the host galaxy performed with the Palomar 200-in telescope four weeks after the burst indicates a redshift of $z$~$=~0.450~\\pm~0.005$ \\citep{djorg01a}. ", + "conclusions": "Table 1 summarizes the magnitudes and upper limits in the various filters from all the above observations. We transformed the fluxes measured in various filters to the $R$ filter by normalizing a $\\beta=-2.3$ power-law spectrum \\citep{price01a,kulkarni01a} to the effective wavelength and flux of each data point, and then plotting the point at the effective wavelength of the $R$-band. This calculated values are listed in the last column of Table 1. Figure~\\ref{fig:lightcurve} shows the resulting light curve of the optical afterglow of GRB~010921. We fit the data with a power law decay plus a constant host galaxy flux, $F = F_0(t-t_0)^{-\\alpha} + F_{host}$. We obtain a best fit decay index of $\\alpha =1.75 \\pm 0.28$ applying a host galaxy magnitude of $R = 21.93$ from the USNO measurements performed 56 days after the burst. This is a typical value for an optical afterglow prior to the jet break, which is predicted to take place $\\sim$ 130 days after the GRB based on the observed energetics of this burst \\citep{djorg01a}. The optical afterglow reported by \\citet{price01a} was $R = 19.6 \\pm 0.3$ on September 22 and significantly fain.ter on September 23 (we place a limit on $R$ of roughly 20.5). We attempt to constrain the early-time power law decay by extrapolating the best-fit power-law decay model back to the LOTIS upper limit of $R > 15.0$ on September~21.256. Figure~\\ref{fig:lightcurve} shows that the early-time LOTIS upper limits are inconsistent with an unchanging decay index from $t - t_0 = 52$~min to $t - t_0 > 20$~h. This may suggest that the optical emission peaked at a magnitude fainter than the LOTIS limiting magnitude, (perhaps similar to the afterglow of GRB 970508) or that the slope changed between the observations, as suggested in the case of GRB 991208 \\citep{castro}. Complex light curve shapes at very early times have been observed (e.g., GRB970805) and can be explained in terms of a distribution of Lorentz factors produced by the central engine \\citep{RandM} and the complex evolution of multi-component shock emission in a relativistic fireball \\citep{kobayashi}. Our early time observations show that afterglow behavior can change quickly within hours after the burst. With rapid localizations we can probe the transition from the prompt emission phase to the subsequent unfolding of the canonical afterglow phase. GRB010921 only provided an upper limit 52 minutes after outburst, but HETE-2 triggers should eventually allow us to obtain simultaneous flux measurements. Robotic telescopes like LOTIS and Super-LOTIS are well suited to finding this early-time emission in response to a near real-time localization of the GRB by HETE-2." + }, + "0112/astro-ph0112504_arXiv.txt": { + "abstract": "Data on galaxies at high redshift, identified by the Lyman-break photometric technique, can teach us about how galaxies form and evolve. The stellar masses and other properties of such Lyman break galaxies (LBGs) depend sensitively on the details of star formation. In this paper we consider three different star formation prescriptions, and use semi-analytic methods applied to the now-standard $\\Lambda$CDM theory of hierarchical structure formation to show how these assumptions about star formation affect the predicted masses of the stars in these galaxies and the masses of the dark matter halos that host them. We find that, within the rather large uncertainties, recent estimates of the stellar masses of LBGs from multi-color photometry are consistent with the predictions of all three models. However, the estimated stellar masses are more consistent with the predictions of two of the models in which star formation is accelerated at high redshifts $z\\gsim3$, and of these models the one in which many of the LBGs are merger-driven starbursts is also more consistent with indications that many high redshift galaxies are gas rich. The clustering properties of LBGs have put some constraints on the masses of their host halos, but due to similarities in the halo occupation of the three models we consider and degeneracies between model parameters, current constraints are not yet sufficient to distinguish between realistic models. ", + "introduction": "A great deal of effort devoted to determining the cosmological parameters has recently paid off. But, although there is good evidence that the cosmological parameters are roughly $\\Omega_m=0.3$, $\\Omega_\\Lambda=0.7$, and $h=0.7$, and that $\\Lambda$CDM with these parameters is a good fit to the observed universe \\cite{isss1}, this theory does not make unique predictions regarding the masses and other properties of galaxies at high redshift. Galaxy properties in cosmological theories also depend on assumptions about uncertain aspects of star formation, supernova feedback, and dust obscuration. Here we will focus on star formation. We consider three different models of star formation, differing in the way that the star formation rate depends on galaxy properties, and discuss the implications for masses, clustering, and other properties of Lyman break galaxies (LBGs) in semi-analytic models. These models \\cite{spf} all assume exactly the same underlying $\\Lambda$CDM model with the parameters above, so the properties of the dark matter halos at any given redshift and the halo merging histories are the same. We also make the same assumptions in each model regarding the initial mass function (IMF), which we assume to be Salpeter between 0.1 and 100 $M_\\odot$. We use the GISSEL00 stellar population synthesis models of Bruzual \\& Charlot \\cite{bc} with solar metallicity, and a simple model for dust obscuration, in which the optical depth is a power-law function of the unobscured ultraviolet luminosity: \\begin{equation} \\tau_{UV} = \\tau_{UV*} (L_{UV}/L_{UV*})^\\beta \\; , \\label{eq:dust} \\end{equation} with $\\tau_{UV*}$ an adjustable parameter and $\\beta=0.5$ \\cite{wangheckman}. The modern approach to semi-analytic modeling was pioneered by White \\& Frenk \\cite{wf91}, and further developed by them and their collaborators in \\cite{kwg93} and \\cite{cole94}. In \\cite{sp}, we reviewed and extended this work, and applied it to high-redshift galaxies \\cite{spf} using three simplified models of star formation. The three star formation models we consider span the range of models proposed for high redshift galaxy formation. The simplest model, termed Constant Efficiency Quiescent, assumes that the quiescent star formation rate per unit mass of cold gas is constant: \\begin{equation} \\dot m_* = \\frac{m_{cold}}{\\tau_*} \\qquad\\qquad {\\rm (CEQ)} \\; . \\end{equation} We showed in \\cite{spf} that the resulting predictions of this model are similar to those of \\cite{cole94} and to more detailed treatment \\cite{baugh98,governato98} from the same group that included starbursts from major mergers. An alternative Accelerated Quiescent model assumes that \\begin{equation} \\dot m_* = \\frac{m_{cold}}{\\tau_{dyn}} \\qquad\\qquad {\\rm (AQ)} \\; . \\end{equation} The predictions of our AQ model are like those of \\cite{kwg93} even though those authors included starbursts from major mergers, because the above star formation prescription (based on data on star formation in nearby galaxies \\cite{kennicutt}) converts gas to stars so efficiently at high redshifts (since typically $\\tau_{dyn} \\propto 1/(1+z)$). Finally, we consider a Collisional StarBurst model in which the quiescent star formation efficiency is constant, but where there is also a burst mode of star formation triggered by galaxy interactions: \\begin{equation} \\dot m_* = \\frac{m_{cold}}{\\tau_*} + (\\dot m_*)_{\\rm starbursts} \\quad {\\rm (CSB)} \\; , \\end{equation} where $(\\dot m_*)_{\\rm starbursts}$ is due to bursts in merging galaxies. The efficiency of star formation in these bursts is scaled according to a model based on hydrodynamical simulations \\cite{mihoshernquist}, in which the efficiency scales as a power-law function of the mass ratio of the merger (see \\cite{spf} for details). We find in our CSB semi-analytic model that most of the star formation at redshifts above unity occurs in starbursts driven by minor mergers, in which the merging satellite has mass less than $\\frac{1}{3}$ that of the central galaxy. As noted, our CEQ model is similar to the Durham group model of several years ago \\cite{baugh98}. At this conference, Baugh presented preliminary results from an alternative model that is similar to our CSB model. \\begin{figure} \\noindent \\begin{minipage}[t]{2.25in} \\centering % \\psfig{file=primackF1a.ps,width=2.25in} \\end{minipage} \\hfill \\begin{minipage}[t]{2.25in} \\centering % \\psfig{file=primackF1b.ps,width=2.25in} \\end{minipage} \\caption{Madau plot of star formation rate density (a) and cosmological density of neutral hydrogen (b) for our three models CSB, AQ, and CEQ. (From \\protect\\cite{spf}.)} \\end{figure} Figure 1a shows the star formation rate density as a function of redshift for these three models. The CEQ model does not produce as many stars at high redshifts $z\\gsim3$ as the extinction-corrected observations indicate (see \\cite{spf} for references), while both AQ and CSB models are acceptable in this regard, indicating that star formation at high redshift must be more efficient than locally. This argument is reinforced by the failure of the CEQ model to produce as many stars at $z>2$ as are indicated by the fossil evidence (see Fig. 12 of \\cite{spf}). However, the AQ model converts gas to stars so efficiently that it may not have as much neutral hydrogen at $z>2$ as is indicated by the data on damped Lyman-alpha systems --- see Fig. 1b. The AQ model may also not have enough gas to fuel quasars at high redshifts \\cite{kaufhaehnelt}, while the CSB model seems acceptable. Further evidence that favors the CSB over the AQ model comes from the predicted LBG luminosity function. The CSB model predicts as many bright LBGs as are observed, although it slightly overpredicts the number of fainter ones, possibly because the dust obscuration prescription Eq. (\\ref{eq:dust}) is unrealistic in predicting very little extinction for lower-luminosity galaxies. But the AQ luminosity function predicts fewer bright LBGs than observed at $z=3$, and far fewer than observed at $z=4$ (see Figs. 4-7 of \\cite{spf}). Bright LBGs only occur in massive halos in the CEQ and AQ models, and there are fewer such halos at higher redshifts. ", + "conclusions": "While the masses of LBGs would seem to provide key information about their nature, due to the considerable uncertainties in our modeling and in deriving stellar or halo masses from the data, it is not currently possible to rule out any of the three very different recipes for star formation considered here on this basis. All three models are also roughly consistent with recent observational estimates of LBG clustering. However, the CEQ model predicts systematically higher stellar masses and also far too few bright LBGs especially at higher redshifts, and the AQ model may use up gas too efficiently to be consistent with other data. One way to check whether it is really true that many of the high-redshift bright galaxies are collision-driven starbursts is to see whether the morphologies of these objects resemble those produced in hydrodynamical simulations of interacting gas-rich galaxies, which are presently underway (see \\cite{racheldisks} for preliminary results)." + }, + "0112/astro-ph0112218_arXiv.txt": { + "abstract": "We examine linear and quasi-liner stages of Cherenkov-drift instability in the relativistic magnetized electron-positron plasma. The external magnetic field lines are assumed to be slightly curved. In this case the curvature drift of relativistic beam particles plays decisive role in the development of the instability. Quasi-linear relaxation of the relativistic beam leads to diffusion of the resonant particles in the momenta space. The expressions for diffusion coefficients of Cherenkov-drift instability are obtained. ", + "introduction": "Introduction} Cherenkov-drift instability was suggested by Kazbegi, Machabeli \\& Melikidze \\cite{kmm89,kmm91c,kmm92b,kmms96} as a possible mechanism for generation of pulsar radio emission and later it was approved in \\cite{lmb99}. In those works the linear theory of Cherenkov-drift instability was developed. It was shown, that in the pulsar magnetosphere due to Cherenkov-drift instability the orthogonally polarized plasma waves are excited. These waves can escape from the magnetosphere and reach observer as a pulsar radio emission. The necessary condition for the development of Cherenkov-drift instability (as for an usual Cherenkov instability) is a presence of a beam of particles in the relativistic magnetized pair plasma (consisted of relativistic electrons $e^{-}$ and positrons $e^{+}$). Generally, Cherenkov type instabilities develop due to a resonant interaction between waves and particles of a beam. The resonance occurs, when the electric field vector ${\\mathbf E}$ and the wave vector ${\\mathbf k}$ of generated waves have got components along direction of the beam velocity ${\\mathbf v}$ (${\\mathbf E}\\cdot{\\mathbf v}\\neq 0$ and ${\\mathbf k}\\cdot{\\mathbf v}\\neq 0$). So, in the magnetized plasma the transverse waves (${\\mathbf E}\\perp{\\mathbf B}\\perp{\\mathbf k}$) propagating along the external magnetic field (~${\\mathbf B}_0\\parallel {\\mathbf k}$~) cannot be generated by an usual Cherenkov instability (because it develops on the beam particles moving along the external straight magnetic field lines: ${\\mathbf v}\\parallel{\\mathbf B}_0$ and ${\\mathbf E}\\cdot{\\mathbf v}= 0$ ). Cherenkov-drift instability develops when the beam particles move along slightly curved magnetic field (SCMF) lines and, hence, drift across the plane where the curved lines lie. The drift motion of beam particles provokes generation of purely transverse as well as longitudinal-transverse waves. Generally there are two most important effects caused by the the particle relativistic motion along the SCMF line: curvature drift and curvature radiation. The drift velocity is directed across the plane of the SCMF lines and is given by the following expression: \\begin{equation} u_d=\\frac{\\gamma v_{\\parallel}^2}{\\omega_B R_B}. \\label{udr} \\end{equation} Here $u_d$ denotes the drift velocity of electrons (positrons drift with the same velocity but in the opposite direction); $\\omega_B= eB/mc$ is the cyclotron frequency of electrons; $R_{B}$ is the curvature radius of the magnetic field line; $\\gamma$ is Lorentz factor of a particle; $c$ is speed of light and $v_{\\parallel}$ is the component of ${\\mathbf v}$ along the magnetic field line. If $\\gamma\\gg 1$, the value of drift velocity $u_d$ could be significant. A single particle, moving along the curved field line, radiates so called curvature radiation which can be easily described as a synchrotron radiation in an effective magnetic field (see e.g. \\cite{z96}). In 1975 Blandford \\cite{b75} investigated the curvature radiation of plasma flowing along the SCMF lines. The problem was studied in the limit of infinite magnetic field $B_0\\rightarrow\\infty$ and it was shown that there is no radiation at all: the waves, radiated by each particle, are absorbed by another one. This result was confirmed later in papers \\cite{zs79,m80,cs88,lm92,kmm91c}, where spatially unlimited plasma flow was considered. Then in the paper by Asseo, Pellat and Sol \\cite{aps83} the sharp boundary was assumed at the edge of the flow propagating along the curved field line and the possibility of the waves excitation was shown at this boundary. If plasma flow has zero width the instability is reduced to that of Goldreich-Keeley \\cite{gk71}. Development of Cherenkov type instability, taking into account the curvature drift motion, was studied and growth rate was calculated in \\cite{kmm86,kmm89,kmm91c,kmms91,kmm92a,kmm92b,pmt92,pmt94} for different particular cases. These results were confirmed later after thorough investigation of the problem in \\cite{lmb99,lbm99}. The instability was called a Cherenkov-drift instability. Presence of the curved magnetic field lines is the necessary condition for both the Cherenkov-drift radiation in plasma and the single particle curvature radiation in vacuum. However, the Cherenkov-drift radiation still could not be interpreted as a plasma curvature radiation, analogous to the single particle curvature radiation: as it will be shown below, the Cherenkov-drift radiation is not generated in the case of $u_{d}\\rightarrow 0$. On the other hand, it is evident that the drift velocity equals to $u_{d}\\approx 0$ if we assume that $B_0\\rightarrow\\infty$. However a single particle radiates even for the infinite intensity of magnetic field. Moreover, the single particle radiates the vacuum wave, while the proper waves of the medium (i.e. relativistic electron-positron plasma) are generated in the case of Cherenkov-drift instability. This particular point was not considered in the works by Blandford \\cite{b75} and Melrose \\cite{m80}. Polarization of these waves strongly differs from that of vacuum waves. Brief examination of the linear theory of Cherenkov-drift instability is discussed in section 2. In section 3 the quasilinear equations for the Cherenkov-drift instability are obtained. In section 4 coefficients describing diffusion of particles in momentum space are evaluated. Alteration of plasma distribution function is studied. The results are summarized in section 5. ", + "conclusions": "In this paper we discuss the development of Cherenkov-drift instability in relativistic magnetized pair plasma. We are taking into account particle drift motion across the plane of SCMF lines, which is significant for the particles of relativistic beam penetrating the bulk pair plasma. We studied quasilinear stage of the instability -- quasilinear interaction of excited waves over plasma particles and corresponding redistribution of particle momenta. As a result, diffusion of particle momenta takes place along, as well as across, the magnetic field lines. The linear stage of Cherenkov-drift instability develops similarly to the usual Cherenkov instability. The mechanism of wave excitation is based on the well known Cherenkov wave-particle interaction. Same as in the case of Cherenkov instability, presence of high energy particle beam (with positive slope on the shape of distribution function) is necessary condition for developing of Cherenkov-drift instability. However, in the case of Cherenkov-drift instability, the generation of oscillations with ${\\mathbf E} \\perp {\\mathbf B}_0$ is possible only due to particle drift motion across ${\\mathbf B}_0$ (contrary to the case of usual Cherenkov instability generating only longitudinal oscillations with ${\\mathbf E}\\parallel {\\mathbf k} \\parallel {\\mathbf B}_0$). Cherenkov-drift instability generates both $t$ (with $E_{\\perp}$) and $lt$ (with $E_{\\perp}$ and $E_{\\parallel}$) waves; electromagnetic oscillations with $E_{\\parallel}$ are generated by particle longitudinal motion with velocities $v_{\\parallel}$, as in the case of usual Cherenkov instability. Back reaction of excited waves over resonant particles should suppress the reason of wave excitation: in the case of Cherenkov-drift instability, the process causes, at the same time, formation of plateau on the distribution function of parallel momenta (similar to the quasilinear case of usual Cherenkov instability) and energy transfer from parallel motion of particles to their motion across the magnetic field. The later process, inhibiting anisotropy in momentum space, is similar to the quasilinear relaxation of cyclotron instability. The reason for development of cyclotron instability -- anisotropy in momentum space ($p_{\\perp}\\ll p_{\\parallel}$) -- is suppressed by particle diffusion over perpendicular momenta. Perpendicular diffusion is described by nonzero diffusion coefficients $D_{\\perp\\perp}$ and $D_{\\perp\\parallel}$. As a result, the energy of parallel motion of beam particles is transferring into the perpendicular energy until $p_{\\perp}\\sim p_{\\parallel}$. As for Cherenkov-drift instability, relaxation is saturated since $p_{\\perp}\\sim p_{d}$ and the rate of alteration of distribution function over perpendicular momenta is described by nonzero perpendicular diffusion coefficients $D_{\\perp\\perp}$ (\\ref{Dtper}) and (\\ref{Dltper}). It is worth to note that, in the Cherenkov-drift instability, the plateau on the parallel distribution function forms not only due to parallel diffusion of particles (which transfers energy from high velocity resonant particles to those with low velocities as for usual Cherenkov instability), but also because of perpendicular diffusion which transfers energy from parallel to perpendicular motion of particles. This scenario works if the other factors which can balance quasilinear diffusion, are not taken into account. Such factors could be, on the one hand, the radiation reaction force (acting on synchrotron emitting particle, spiraling in strong magnetic field) and, on the other hand, the force arising due to particle motion in weekly inhomogeneous field. We plan to include these factors into consideration in the future works." + }, + "0112/astro-ph0112243_arXiv.txt": { + "abstract": "GAIA is the ``super-Hipparcos'' survey satellite selected as a Cornerstone 6 mission by the European Space Agency. GAIA can measure microlensing by the brightening of source stars. For the broad G band photometer, the all-sky source-averaged photometric optical depth is $\\sim 10^{-7}$. There are $\\sim 1300$ photometric microlensing events for which GAIA will measure at least one datapoint on the amplified lightcurve. GAIA can also measure microlensing by the small excursions of the light centroid that occur during events. The all-sky source-averaged astrometric microlensing optical depth is $\\sim 2.5 \\times 10^{-5}$. Some $\\sim 25000$ sources will have a significant variation of the centroid shift, together with a closest approach, during the lifetime of the mission. This is not the actual number of events that can be extracted from the GAIA dataset, as the false detection rate has not been assessed. A covariance analysis is used to study the propagation of errors and the estimation of parameters from realistic sampling of the GAIA datastream of transits in the along-scan direction during microlensing events. The mass of the lens can be calculated to good accuracy if the lens is nearby so that angular Einstein radius $\\thetaE$ is large; if the Einstein radius projected onto the observer plane $\\rE$ is about an astronomical unit; if the duration of the astrometric event is long ($\\gta 1$ year) or if the source star is bright. Monte Carlo simulations are used to study the $\\sim 2500$ events for which the mass can be recovered with an error of $< 50 \\%$. These high quality events are dominated by disk lenses within a few tens of parsecs and source stars within a few hundred parsecs. We show that the local mass function can be recovered from the high quality sample to good accuracy. GAIA is the first instrument with the capabilities of measuring the mass locally in very faint objects like black holes and very cool white and brown dwarfs. For only $\\sim 5\\%$ of all astrometric events will GAIA record even one photometric datapoint. There is a need for a dedicated telescope that densely samples the Galactic Centre and spiral arms, as this can improve the accuracy of parameter estimation by a factor of $\\sim 10$. The total number of sources that have an astrometric microlensing variation exceeding the mission target accuracy is $\\sim 10^5$. The positional measurement of one source in every twenty thousand is affected by microlensing noise at any instant. We show that microlensing is negligible as an unbiased random error source for GAIA. ", + "introduction": "GAIA is the European Space Agency (ESA) satellite now selected as a Cornerstone 6 mission as part of the Science Program~\\footnote{http://astro.estec.esa.nl/gaia}. It is a survey satellite that provides multi-colour, multi-epoch photometry, astrometry and spectroscopy on all objects brighter then $V\\approx 20$ (e.g., ESA 2000; Perryman et al. 2001). The dataset is huge with information of unprecedented precision on over a billion objects in our Galaxy alone. GAIA is the successor to the pioneering {\\it Hipparcos} satellite, which flew from 1989 to 1993. The {\\it Hipparcos} program was both simpler and smaller: it measured only $10^5$ rather than $10^9$ objects and it was provided with a target list rather than being left to determine its targets for itself, as GAIA will. The {\\it Hipparcos} final results were released only when the mission was complete, whereas many of the science goals of GAIA, especially for bursting or time-varying phenomena like microlensing, may require an early analysis and release of preliminary data. GAIA carries out continuous scanning of the sky. The satellite rotates slowly on its spin axis, which itself precesses at a fixed angle to the Sun of $55^\\circ$. As GAIA rotates, light enters the entrance chambers, reflects off mirrors and falls on the focal planes of one of three telescopes, two of which (ASTRO-1, ASTRO-2) measure the positions of stars, one of which (SPECTRO) performs spectroscopy. GAIA observes in three directions along a great circle simultaneously. The astrometric parameters of stars are recovered from the time series of the one-dimensional transits distributed over the five year mission lifetime. A small fraction of the objects monitored by GAIA will show evidence of microlensing. GAIA can observe microlensing by measuring the photometric amplification of a source star when a lens and a source are aligned. This is the approach followed by the large ground-based microlensing surveys like MACHO, EROS, OGLE and POINT-AGAPE (see e.g., Alcock et al. 1997; Aubourg et al. 1995; Udalski et al. 1994, Auri\\`ere et al. 2001). GAIA is inefficient at discovering photometric microlensing events, as the sampling of individual objects is relatively sparse (there are a cluster of observations once every two months on average). \\begin{figure} \\epsfxsize=8cm \\centerline{\\epsfbox{event0.ps}} \\caption{This shows the relative right ascension and declination of a source with (solid line) and without (dashed line) a microlensing event. The trajectory is shown over the GAIA mission lifetime of 5 years. Note that the deviations caused by microlensing are present well before and after the time of maximum of the event. (The lens is at 150 pc, the source at 1.5 kpc, the transverse velocity is $70 \\kms$ and the impact parameter $u$ is 1.5. The lens has mass $0.5\\msun$.)} \\label{fig:pathofsource} \\end{figure} However, there is a much more powerful strategy available to GAIA. Although the two images of a microlensed source are unresolvable, GAIA can measure the small deviation (of the order of a fraction of a milliarcsec) of the centroid of the two images around the trajectory of the source. Astrometric microlensing is the name given to this excursion of the image centroid. The cross section of a lens is proportional to the area it sweeps out on the sky, and so to the product of lens proper motion and angular Einstein radius. Each of these varies as the inverse square root of lens distance, so the signal is dominated by nearby lenses. The detection of astrometric microlensing events for pointed observations has been considered many times before (e.g., Walker 1995; Miralda-Escud\\'e 1996; Paczy\\'nski 1996; Boden, Shao \\& van Buren 1998). Once alerted by a ground-based survey, the events are followed by narrow angle differential astrometry. The detection of astrometric microlensing using observations from a scanning satellite has been considered before by H{\\o}g, Novikov \\& Polnarev (1995) in the context of the proposed ROEMER mission (a forerunner of GAIA). We will show that the all-sky source-averaged astrometric microlensing optical depth is $\\sim 2.5 \\times 10^{-5}$, over an order of magnitude greater than the photometric microlensing optical depth. There are two main difficulties facing GAIA in exploiting this comparatively high probability. First, the astrometric accuracy of a single measurement by GAIA depends on the source magnitude and quickly degrades at magnitudes fainter than $G \\approx 15$. Second, GAIA provides a time-series of one-dimensional astrometry. The observed quantity is the CCD transit time for the coordinate along the scan. This is the same way the Hipparcos satellite worked (see {\\it The Hipparcos and Tycho Catalogues}, ESA (1997), volume 3, section 16). From the sequence of these one-dimensional measurements, the astrometric path of the source, together with any additional deflection caused by microlensing, must be recovered. This paper assesses the microlensing signal that will be seen by GAIA. Section 2 gives the formulae for the photometric and astrometric microlensing optical depths. These are used to estimate the total number of microlensing events that GAIA can measure. In Section 3, a covariance analysis is used to demonstrate the effects of error propagation on the recovery of the parameters of microlensing events. For the purposes of GAIA, we show that the disk stars within a few hundred parsecs of the Sun are the most important source and lens populations. Section 4 describes the extinction law and the Galaxy model used to generate microlensing events for our Monte Carlo simulations. The synthetic data are sampled with GAIA's scanning law and realistic errors are applied to provide the one-dimensional astrometric datastreams. Section 5 discusses the results of the simulations, both for the entire sample of events and for the subset of gold-plated events whose parameters can be recovered to good accuracy. Finally, Section 6 examines the overall strategy for identification of events, which need to be distinguished from other forms of astrometric deviation such as binary companions. \\begin{figure} \\epsfxsize=8cm \\centerline{\\epsfbox{event1.ps}} \\caption{Upper panel: Astrometric shift of the microlensing event of Figure 1, as seen by a barycentric (dotted line) and a terrestrial observer (solid line). Lower panel: Simulated data incorporating typical sampling and astrometric errors for GAIA. Also shown for comparison are the theoretical trajectories of the source with (grey line) and without (dashed line) the event. The insets show the deviations at the beginning and the midpoint of this high signal-to-noise event. (The accuracy $\\sigmaa$ of the astrometry is $300\\muas$, corresponding roughly to a $17$th magnitude star).} \\label{fig:examplesone} \\end{figure} \\begin{figure} \\epsfxsize=8cm \\centerline{\\epsfbox{event2.ps}} \\caption{Same as Figure 2, but the lens distance is 1 kpc. This shows a microlensing event with a $5 \\sqrt{2} \\sigmaa$ variation of the centroid shift. GAIA will be able to measure the astrometric deviation but will not be able to extract any useful information on the microlensing parameters.} \\label{fig:examplestwo} \\end{figure} \\begin{table*} \\begin{center} \\begin{tabular}{c|cccccccccccc}\\hline $G$ (in mag) & 10 & 11 & 12 & 13 & 14 & 15 & 16 & 17 & 18 & 19 & 20 \\\\ $\\sigmap$ (in mmag) & 9 & 9 & 9 & 9 & 9 & 10 & 12 & 16 & 23 & 36 & 60 \\\\ $\\sigmaa$ (in $\\mu$as) & 30 & 30 & 30 & 40 & 60 & 90 & 150 & 230 & 390 & 700 & 1400 \\\\ \\hline \\end{tabular} \\end{center} \\caption{This table lists the mean accuracy in photometry $\\sigmap$ and in position $\\sigmaa$ versus $G$ band magnitude for the GAIA satellite. Note that $\\sigmaa$ is the accuracy of a single astrometric measurement, not the target accuracy at the end of the GAIA mission, which is better by a factor of $\\sim 10$. The values are approximate sky averages, adapted from Tables 7.3 and 8.2 of ESA (2000).} \\label{table:astroacc} \\end{table*} \\begin{figure} \\epsfysize=15.cm \\centerline{\\epsfbox{depth_all.ps}} \\caption{Upper panel: An all-sky map of the source-averaged astrometric optical depth. Meridians of galactic latitude are shown at $60^\\circ$ intervals, parallels of longitude at $30^\\circ$ intervals. The contour levels are marked. The number within each contour refers to the mean value of the optical depth inside the shaded region enclosed between the contours. All numbers are in units of $10^{-5}$. Middle panel: As the upper panel, but for the photometric microlensing optical depth. All numbers are in units of $10^{-7}$. Lower panel: An all-sky map of the starcounts. The number within each contour refers to the total number of stars inside the shaded region enclosed between the contours. To find the instantaneous number of events in any region, we multiply the number of stars by the average optical depth. (For details of the source and lens populations and extinction model, see Section 4).} \\label{fig:maps} \\end{figure} ", + "conclusions": "GAIA is an ambitious astrometric survey satellite that is planned for launch no later than 2012 by the European Space Agency (see ESA 2000; Perryman et al. 2001). GAIA scans the whole sky, performing multi-colour and multi-epoch photometry and spectroscopy, as well as astrometry. The reference frame is tied to a global astrometric reference frame defined by extragalactic objects. For bright sources, GAIA has the capabilities to perform individual astrometric measurements with microarcsecond accuracy. Even for objects as faint as $V \\approx 20$, the positions, proper motions and parallaxes will be recovered to within at worst a few hundred microarcsecs at the end of the mission. GAIA can measure microlensing by the small excursions of the light centroid that occur during microlensing. We use a stringent definition of an astrometric microlensing event as one with a significant variation ($5 \\sqrt{2} \\sigmaa$) of the centroid shift, together with a closest approach during the lifetime of the GAIA mission. The all-sky averaged astrometric microlensing optical depth is $\\sim 2.5 \\times 10^{-5}$. This means that $\\sim 25000$ sources will exhibit astrometric microlensing events with such characteristics during the course of the mission. GAIA can also measure microlensing by the photometric brightening that accompanies low impact parameter events. The all-sky averaged photometric optical depth is $\\sim 10^{-7}$, so there are $\\sim 3600$ photometric microlensing events during the five year mission lifetime, most of which are undetectable because of the poor sampling. Consequently, very few astrometric events are measured photometrically as well. In fact, only $\\sim 1260$ astrometric events have at least one datapoint sampled on the photometric lightcurve. Only for $\\sim 427$ of the astrometric events will GAIA itself obtain enough photometric datapoints to improve substantially the characterisation of the event. Let us again emphasise that these numbers only refer to numbers of microlensing events with specified characteristics. Any statement as to the actual number of microlensing events that GAIA will find must take into account the number of false detections thrown up by the identification algorithm. Our event statistics refer only to the signal provided by stellar lenses in the Galactic disk. We have not taken into account either stellar lenses in the Galactic bulge or dark objects in the halo (popularly called Machos). This is reasonable, as the cross-section for astrometric microlensing favours nearby lenses. The most valuable events are those for which the Einstein crossing time $\\tE$, the angular Einstein radius $\\thetaE$ and the relative parallax of the source with respect to the lens $\\pi_{\\rm sl}$ can all be inferred from GAIA's datastream. The mass of the lens then follows directly. If the source distance is known -- for example, if GAIA itself measures the source parallax -- then a complete solution of the microlensing parameters is available. Of these quantities, it is the relative parallax that is the hardest to obtain accurately. A covariance analysis is used to follow the propagation of errors and establish the conditions for recovery of the relative parallax. This happens if the angular Einstein radius $\\thetaE$ is large and the Einstein radius projected onto the observer plane $\\rE \\sim 1$ au so that the Earth's motion about the Sun gives a substantial distortion. It is also aided if the source is bright so that GAIA's astrometric accuracy is high and if the duration of the astrometric event is long so that GAIA has time to sample it fully. These conditions favour still further lensing populations that are very close. Monte Carlo simulations are used to establish the characteristics of the $\\sim 10 \\%$ of events for which GAIA can recover the mass of the lens to good accuracy. The typical lens distance is $\\sim 50$ pc and the typical source distance is $\\sim 300$ pc. The highest quality events seen by GAIA are overwhelmingly dominated by very local lenses. We conclude that one of the major scientific contributions of microlensing studies with GAIA will be the determination of the mass function in the solar neighbourhood. Of course, direct mass measurements are presently possible just for binary stars with well-determined orbits. Microlensing is the only technique which can measure the masses of individual stars. GAIA is the first instrument with the ability to survey the astrometric microlensing signal provided by nearby lenses. We have used Monte Carlo simulations to show that GAIA can reconstruct the mass function in the solar neighbourhood from the sample of its highest quality events. This works particularly well for masses exceeding $\\sim 0.3 \\msun$. Below $0.3 \\msun$, the reconstructed mass function tends to underestimate the numbers of objects, as the highest quality events are biased towards larger angular Einstein radii. If there are local populations of low mass black holes, or very cool halo and disk white dwarfs or very old brown dwarfs, then they will have easily eluded detection with available technology. However, the astrometric microlensing signal seen by GAIA will be sensitive to local populations of even the dimmest of these stars and the darkest of these objects. GAIA is the first instrument that has the potential to map out and survey our darkest neighbours." + }, + "0112/astro-ph0112075_arXiv.txt": { + "abstract": "We present new VLT spectroscopic observations of the most distant quasar known, SDSS J1030+0524 at z$=$6.28, which was recently discovered by the Sloan Digital Sky Survey. We confirm the presence of a complete Gunn-Peterson trough caused by neutral hydrogen in the intergalactic medium. There is no detectable flux over the wavelength range from 8450 to 8710 \\AA. We set a stronger limit on the drop of the flux level blueward of the Ly$\\alpha$ line: a factor of $>$ 200. Below 8450 \\AA \\ the spectrum shows a rise in flux, with a large fraction ($>$ 60 $\\%$) of the total emission produced by few narrow features of transmitted flux. We discuss the proximity effect around this quasar, with the presence of transmitted flux with many absorption features in a region of about 23$h^{-1}$ comoving Mpc. If assuming the surrounding medium were completely neutral, the size of this region would imply a quasar lifetime of $\\sim$ 1.3 $\\times$10$^7$ years. We also present near-IR spectroscopy of both SDSS J1030+0524 and of SDSS J1306+05, the second most distant quasar known, at redshift 6.0. We combine measurements of the CIV line and limits on the HeII emission from the near-IR spectra with the NV line measurements from the optical spectra to derive the metal abundances of these early quasar environments. The results are indistinguishable from those of lower redshift quasars and indicate little or no evolution in the metal abundances from z$\\sim$ 6 to z$\\sim$2. The line ratios suggest supersolar metallicities, implying that the first stars around the quasars must have formed at least a few hundreds of Myrs prior to the observation, i.e., at redshifts higher than 8. ", + "introduction": "Quasars are amongst the most luminous objects in the Universe, allowing us to study them and any intervening material out to very large distances, corresponding to look-back times when the Universe was very young. Hence finding and studying quasars at high redshifts is one of the best ways to constrain the physical conditions in the early Universe. The mere existence of luminous quasars at such early times, and the implied presence of black holes with M$\\ \\ge 10^9$ M$_\\odot$ place stringent limits on the epoch at which massive condensed structures formed, thereby constraining structure formation models (e.g.\\ Efstathiou and Rees 1988). At high redshifts such luminous quasars must be associated with very massive halos, hence they are expected to be found near high peaks (4-5 $\\sigma$ or more) of the large scale density field (e.g.\\ Efstathiou \\& Rees 1988, Nusser \\& Silk 1993) that collapsed sufficiently early. Furthermore, several arguments suggest that the formation of elliptical galaxies and massive bulges are paralleled by an early quasar phase (e.g.\\ Kauffmann \\& Haehnelt 2000). The spectra of high redshift quasars contain important information on the enrichment history of the gas in the quasar environment, and probe the star formation preceding the epoch at which the quasars are observed, possibly the first stars that formed in massive collapsed structures (e.g.\\ Hamann \\& Ferland 1999). Finally, high redshift quasars serve as probes of the intergalactic medium via the absorption of the Lyman $\\alpha$ forest. The Sloan Digital Sky Survey (SDSS -- York et al. 2000) has, amongst its scientific aims, the construction of the largest sample of quasars ever, with more than 10$^5$ objects spanning a large range of redshift and luminosities. The SDSS has already found an unprecedented number of new high redshift quasars, including more than 200 new quasars at z$\\ge$4 (e.g. Fan \\etal\\ 2000, Zheng \\etal\\ 2000, Anderson \\etal\\ 2001, Fan et al. 2001a, Schneider \\etal\\ 2001). These high redshift quasars have been efficiently selected by their distinctive position in color-color diagrams, with characteristic colors due to the main feature of the quasar spectra, viz., the strong Ly$\\alpha$ emission line, the Ly$\\alpha$ forest and the Lyman limit. Recently, Fan et al.\\ (2001b) presented the discovery of SDSS J1030+0524, found during follow-up spectroscopy of $i$-band drop-out objects, i.e. objects showing very red $i^* - z^*$ color and relatively blue $z^*$-J colors (following the notation of previous SDSS papers, we use the superscript $^*$ for the photometry, and the letters alone for the filters, Stoughton \\etal\\ 2002). Moderate signal to noise ratio (S/N) optical spectra taken with the ARC 3.5m telescope showed very strong (rest-frame equivalent width EW$\\sim$50 \\AA) Ly$\\alpha$ + NV emission at $\\sim$8850 \\AA \\ and $\\sim$9050 \\AA, \\ respectively. Based on a fit of the CIV line the redshift was estimated estimated to be z$\\sim$ 6.28 which makes it unambiguously the most distant quasar known. Follow up optical spectroscopic observations with Keck (Becker et al. 2001) have shown the first clear detection of a complete Gunn-Peterson trough (Shklovsky 1964, Scheuer 1965, Gunn \\& Peterson 1965) in the spectrum of this quasar. The flux level drops by a factor of at least 140 relative to an estimate of the unabsorbed continuum level, and the spectrum is consistent with zero flux in the Ly$\\alpha$ forest region immediately blueward of the Ly$\\alpha$ emission line. Even if the existence of the trough itself does not imply that the quasar is observed prior to the re-ionization epoch (e.g.\\ Barkana 2001), the fast evolution of the mean absorption in the spectra of quasars at z$>$ 5 suggests that the universe is approaching the reionization epoch at z$\\sim$ 6 (Becker \\etal\\ 2001, Fan \\etal\\ 2001b, Djorgovski \\etal\\ 2001). In a companion paper (Fan et al. 2001c), we use cosmological simulations to estimate the evolution of the ionizing background and to constrain the redshift of reionization (see also Gnedin 2000, Cen \\& McDonald 2001, Lidz et al. 2001). We have obtained an optical VLT spectrum of the quasar to independently measure the Gunn-Peterson trough and to set tighter limits on the transmitted flux. We have also obtained a near-IR spectrum, to set constraints on the metallicity of the quasar environment and to better estimate the redshift from the CIV emission line, which unlike \\lya\\ is not affected by intervening absorption. Finally we obtained a near-IR VLT spectrum of the second most distant quasar SDSS J1306+0356 at z=5.99 (Fan \\etal\\ 2001b). We will use the results from this spectrum in the discussion of the metallicity in high redshift quasars. The paper is organized as follows: in Section 2 we describe the observations and present the optical and near-IR spectra of the two quasars. In Section 3 we discuss the main results obtained from the optical spectrum: the Gunn-Peterson trough and the limits we can set from the new data, as well as the region in which flux is present, and the proximity effect. In Section 4 we derive the metallicity in the two highest redshift quasars and discuss its implication for the star formation timescales. \\\\ Throughout the paper we assume a flat, $\\Lambda$-dominated universe with $H_{0} = 65$ km s$^{-1}$ Mpc$^{-1}$, $\\Lambda = 0.65$ and $\\Omega = 0.35$ \\cite{OS95,KT95}. ", + "conclusions": "We have presented new optical and near-IR spectroscopic observations of the two most distant quasars known, at redshifts 6.28 and 6.0. We have confirmed the presence of a complete GP trough in the highest redshift object, derived new limits on the transmitted flux and optical depth in this region, and analyzed the characteristics of the flux emission in the region. In Fan et al. (2001c) we use semi-analytic models to study the evolution of the ionizing background and the epoch of reionization from the z$\\sim$ 6 quasars. From the optical and near-IR spectra, we have estimated the metallicities of these early quasar environments, finding supersolar values for both objects. These high metallicities imply that the first stars around the quasars must have formed at least a few hundreds of Mpc prior to the observation, i.e. at redshifts higher than 8. Follow-up observations of other metal emission lines will help us to better constrain the ages of the star forming regions." + }, + "0112/astro-ph0112133_arXiv.txt": { + "abstract": "We present preliminary results of numerical simulations of dissipationless merging of stellar systems, aimed at exploring the consequences of merging between gas free, spheroidal systems. In particular, we study the dynamical and structural characteristics of hierarchical merging between equal mass stellar systems, and we compare the properties of the end-products with the most important structural and dynamical scaling relations obeyed by spheroids. In the explored hierarchy of four successive mergings we find that the FP tilt is marginally conserved, but both the Faber-Jackson and Kormendy relations are {\\it not} conserved. ", + "introduction": "From a {\\it theoretical} point of view, in the scenario of hierarchical galaxy formation elliptical galaxies (Es) formed by merging of smaller systems~\\cite{4,11,12}. On the other hand, from {\\it observations} we know that Es satisfy many tight scaling relations: for example the Fundamental Plane~\\cite{5,6} (FP), the $\\mbh-\\sg0$~\\cite{8,9}, the Mg$_2-\\sg0$~\\cite{1}, and the color--magnitude~\\cite{2} relations. In particular, the FP of Es relates their central velocity dispersion $\\sg0^2$, total luminosity $\\Lb$ and circularized effective radius $\\cRe$, with a 1-sigma scatter of $\\simeq15\\%$ in $\\cRe$ for fixed $\\Lb$ and $\\sg0$~\\cite{10}. {\\it Here we focus on the constraints imposed by the existence of the FP on the role of dissipationless merging in the formation of Es. In other words we want to verify, by using numerical simulations, whether the FP is ``closed'' with respect to the merging process}. Among the motivations of this exploration is the fact that, in the merging of two galaxies with masses $(M_1 ,M_2)$ and virial velocity dispersions $(\\sigma_{v,1},\\sigma_{v,2})$, the virial velocity dispersion of the merger (in case of no mass loss and negligible initial interaction energy of the galaxy pair when compared to their internal energies) is given by \\begin{equation} {\\sigma^2_{v,1+2}}={M_1{\\sigma^2_{v,1}}+M_2{\\sigma^2_{v,2}}\\over{M_1+M_2}}. \\end{equation} It follows that $\\sigma_{v,1+2}\\leq\\,{\\rm max}(\\sigma_{v,1},\\sigma_{v,2})$, i.e., {\\it the virial velocity dispersion cannot increase in a merging process of the kind described above}. On the other hand, the Faber-Jackson relation~\\cite{7} (FJ) indicates that the {\\it projected central velocity dispersion} increases with galaxy luminosity. In addition, since in a purely gas free merging process the stellar mass--to--light ratio cannot increase, {\\it the FP can be maintained only by structural and/or dynamical non--homology}. ", + "conclusions": "" + }, + "0112/astro-ph0112419_arXiv.txt": { + "abstract": " ", + "introduction": "Recent measurements of the luminosity distance $D_L(z)$ using Type Ia supernovae \\cite{{Schmidt:1998ys},{Riess:1998cb},{Perlmutter:1999np}} suggest that an accurate value of $D_L(z)$ may be obtained in the near future. In particular, SNAP \\cite{snap} should provide us the luminosity distances of $\\sim$2000 Type Ia supernovae with an accuracy of a few percent up to $z\\sim1.7$ every year. Also, from observation of the first Doppler peak of the anisotropy of the CMB, it is suggested that the universe is flat, \\cite{ber00,lange00} and this may be proved in the future from obervations by MAP and Planck. Under the assumption of the homogeneity and isotropy of our universe, these observations suggest that dark energy is dominant at present. In an attempt to determine the nature of dark energy, many arguments have been given. \\cite{Weinberg:2000yb} Recently, some mechanisms to account for the observed tiny but finite dark energy are proposed. \\cite{Yokoyama:2001ez,Yokoyama:2002ts} However, at present we do not have a firm and reliable theoretical basis to investigate such a small energy scale compared with the Planck scale. In short, the nature of dark energy under the assumption of the homogeneity and isotropy of our universe is still a great mystery. From the observed isotropy of the CMB, assuming that we are not in a special part of our universe, the universe should be homogeneous. However, if our position in the universe is special, the universe might be inhomogeneous, although the CMB is isotropic. Such cosmological models have been constructed using spherically symmetric models in which we are near the symmetric center. Some authors have considered such models to interpret the SNIa data for small $z$, \\cite{Celerier:2000hp} as well as large $z$ assuming a void structure \\cite{{Tomita:2000a},{Tomita:2001a},{Tomita:2001b}} to avoid dark energy. Such possibilities may be regarded as absurd. However, our point of view in this paper is to construct a possible inhomogeneous dust universe derived from the observed $D_L(z)$. If such a model is consistent with present observational results, the inhomogeneous universe should be examined more seriously, because the dark energy solution is also absurd in the sense that it is $\\sim 120$ orders of magnitude smaller than the Planck scale. In short, we suppose that the question we need to answer to roughly reduces to the following: Which is more absurd, dark energy or an inhomogeneous universe? In the former case, there is no reliable theory to examine the problem at present, while the latter case can be studied in the frame-work of known theories. We would like to point out that it is not taste but, rather, future observations that will confirm either dark energy or an inhomogeneous universe. The analysis of high redshift supernovae gives us the luminosity distance-redshift relation $D_{L}(z)$ along the observational past null cone up to $z \\sim 1$. \\cite{{Schmidt:1998ys},{Riess:1998cb},{Perlmutter:1999np}} The data fit well with $D_{L}(z)$ in the homogeneous and isotropic universe with $\\Omega_m = 0.3$ and $\\Omega_\\Lambda = 0.7$ given by \\begin{equation} \\label{D_L} D_{L}(z) = \\frac{1}{H_0}(1+z) \\int_0^z \\frac{dz'}{\\sqrt{\\Omega_m (1+z')^3 +\\Omega_\\Lambda}}. \\end{equation} In this paper we assume that $D_{L}(z)$ is given by Eq. (\\ref{D_L}) with $\\Omega_m = 0.3$ and $\\Omega_\\Lambda = 0.7 $ for $z \\lesssim 1$. This is done for the sake of simplicity to make the arguments clearer. In particular, we do not wish to claim that $D_{L}(z)$ with $\\Omega_m = 0.3$ and $\\Omega_\\Lambda = 0.7 $ has been confirmed. While $D_{L}(z)$ for $ 1 \\lesssim z < 1.7$ is not certain even at present and will be obtained in the future, for example, by SNAP. Since the scale factor $a$ obeys \\begin{equation} \\frac{\\ddot{a}}{a}=-\\frac{4\\pi}{3}(\\rho + 3p), \\end{equation} $D_{L}(z)$ with $\\Omega_m = 0.3$ and $\\Omega_\\Lambda = 0.7$ implies that the present universe is accelerating, while for the dust universe ($p=0$), $a$ should be decelerating. Therefore it may be concluded that observations is inconsistent with the inhomogeneous dust model. However, the point is that to determine $D_{L}(z)$, we are observing Type Ia supernova events that occurred at past times in spatial positions separated from us. In the inhomogeneous universe model, the time dependence of $a$ at a point separated from us differs from that of our position, so that we may obtain an apparent accelerating universe even though the dust universe is decelerating locally. Before ending this introduction, we comment on some other works relevant to this paper. The inhomogeneous scenario is not the only alternative to dark energy. Giving up the assumption that the cosmic substratum is composed of perfect fluids, bulk pressures that differ from the kinetic pressure can be allowed. The assumption of an effective anti-friction force leads to a model that has only one dark component (CDM) and is consistent with the CMB and SNIa data. \\cite{Zimdahl:2001zm} Also, there is an approach somewhat related to that presented in this paper (though with different motivation) that has been used. \\cite{Avelino:2001a,Avelino:2001b,Avelino:2001c} ", + "conclusions": "In this paper we have constructed inhomogeneous dust models without dark energy. We find that these models are consistent with the observed $D_L(z)$ up to $z=1$, as from Fig. 1 no difficulties are encountered up to $z\\sim 1$ for any set of parameter values in both the BigBang time inhomogeneity and curvature inhomogeneity cases. For $z > 1$, we have difficulties in our inhomogeneous dust models. Recently, the SNIa at a redshift of $\\sim 1.7$ was found \\cite{{Gilliland:1999ne},{Riess:2001gk}} with rather large uncertainties. However, only a single SNIa at a redshift of $\\sim 1.7$ is not enough to construct an accurate $D_L(z)$, although that result seems to rule out the `grey-dust' hypothesis. In addition, the results of a recent investigation of the effect of gravitational lensing on this SNIa suggests that the grey-dust model may be consistent with the observational data. \\cite{Mortsell:2001,Gunnarsson:2001} If future observations confirm $D_L(z)$ up to $z\\sim 2$ with $\\Omega_m \\sim 0.3$ and $\\Omega_\\Lambda \\sim 0.7$ , it can be concluded that our inhomogeneous dust models are incompatible with the observations and that some form of dark energy is likely to exist. However, if future observations confirm that $D_L(z) $ for $z > 1$ is not consistent with Eq. (\\ref{D_L}), the plausibility of our inhomogeneous dust models should be studied more extensively. In such a case, the first Doppler peak as well as the higher ones will give us another constraints on the inhomogeneous universe models. It may believed that the existing observations for $00.2\\,$GeV). \\item The effective 3D space is homogeneous and isotropic at all times when the temperature of the universe is $T>0.2\\,$GeV. \\end{enumerate} In this context, we analyze the possibility of baryogenesis and find specific conditions under which it is feasible. Our setup contains all the necessary ingredients of baryogenesis \\cite{sakharov}: \\begin{enumerate} \\item Baryon number violation (dimension-6 operator) \\item \\( C \\) and \\( CP \\) violation (Standard Model CKM + dimension-6 operator) - the resulting quantity is much larger than the Jarlskog invariant. \\item Nonequilibrium (spacetime expansion or other unspecified means) \\end{enumerate} The reason for the difficulty of this scenario will be ingredient 1, the baryon number violating operator being too weak in the low temperature setting, independently of ingredients 2 and 3. Since the usual Jarlskog invariant of the SM is not the relevant quantity, given the presence of dimension-6 operators, the reason why this baryogenesis scenario is difficult has nothing to do with traditional wisdom. In particular, we find that dimension-6 operators are in most cases too weak to compete with baryon number-conserving operators which create entropy, for most sets of reactants. Na\\\"\\i vely, one finds this result a little surprising because one would expect \\begin{equation} \\eta_{B} \\equiv \\frac{n_{B}}{s}\\sim \\frac{\\delta_{CPV} Bn} {g_{*S}T^{3}_{*}} \\end{equation} can be easily engineered to be sufficiently large, where \\( \\delta_{CPV} \\) is the dimensionless quantity associated with $CP$ violation, \\( B \\) is the maximal branching ratio into the baryon number violating channel fixed by our model, \\( n \\) is the number density of particles that can be converted into baryon number, \\( g_{*S} \\) is the number of degrees of freedom contributing to the entropy, and \\( T_{*} \\) is the temperature at some fiducial time after which the baryon number is conserved. However, even though \\( \\delta_{CPV} \\) can be naturally large in our scenario (not governed by the Jarlskog invariant), \\(B n/(g_{*S}T_{*}^{3}) \\) cannot in general be made large because as the baryon number violating reactions are creating baryon asymmetry, the baryon number conserving counterparts are creating proportionately large amounts of entropy density. In other words, having the quark-lepton separation go to zero does not generate order 1 branching of the baryon number violating channels. Instead, such channel is always suppressed by a scale much larger than the electroweak scale. We do find an exception to our ``no go'' rule for a special situation in which the universe starts out dominated by the $b$ and $c$ quarks and gluons (which might happen through the peculiarities of reheating after inflation), and the lighter quark $c$ is out of equilibrium while the heavier quark $b$ is in equilibrium, for example due to the anomalously large expansion rate and the anomalously small density of $c$ quarks. In this exceptional scenario, there are two key ingredients compensating the suppression of the dimension 6 operator: 1. Due to the small $|V_{cb}|^2$, the $b$ quark interactions with out of equilibrium $c$ quarks do not generate significant entropy; 2. The integration time for baryogenesis can be made long if the expansion rate is small, due to the nonstandard cosmology of extra dimensions. However, this remote possibility, which requires additional assumptions about special initial conditions and/or more new physics, is unlikely to be realizable, and we do not have a detailed model to demonstrate its full viability. Throughout this paper, we do not assume thermal equilibrium initial conditions, except where explicitly noted. Hence, although the short distance operators are fixed, the nonzero initial particle densities that exist in the universe can be {\\em a priori}\\/ arbitrary combinations of species. Thus, when we consider bounds, we explicitly check all likely combinations of species without any bias to their initial number distributions, except when otherwise noted. As far as the out-of-equilibrium-condition due to the expansion of the universe is concerned, we assume that somehow the Planck scale can be adjusted to dial into the appropriate expansion rate. Of course, in a more completely specified model, this may turn out to be problematic, due to the bounds on moduli overproduction and other light moduli problems. However, our results are robust with respect to how the moduli problems and initial conditions problems can be solved. In fact most of our results do not depend on how the particles are out of equilibrium. This broad insensitivity with respect to cosmological details is one of the key features to our work. The order of our presentation will be as follows. We begin with a discussion of the class of braneworlds and the energy scales and couplings that we are interested in. We then present our constraint on this scenario for baryogenesis with dimension-6 operators in the low temperature regime. Finally, we summarize and conclude. ", + "conclusions": "We have studied a new scenario of baryogenesis involving dimension-6 baryon number-violating operators that would generically arise from integrating out ultraviolet degrees of freedom in the context of the Standard Model. The question to be answered in this scenario was whether a time-dependent dimensionless suppression coefficient for the dimension-6 operator can allow sufficient baryogenesis between the effective temperatures of $30\\,$GeV and $0.2\\,$GeV, if the only fields that are allowed to participate in the short distance physics are those of the Standard Model. The upper bound on the temperature is motivated from the cosmological bound on KK graviton decay while the lower bound is purely for our calculational feasibility (related to the QCD phase transition). The time-dependent suppression coefficient was motivated from the generic possibility that with the Arkani-Hamed-Schmaltz mechanism of suppressing proton decay, the proton decay suppression factor can be time-dependent. Unfortunately, if the fundamental scale setting the size of the unsuppressed dimension-6 operator is forced up to values $M_f>100\\,$TeV by the bound from neutron-antineutron oscillations, the branching ratio of $B$-violating interactions to $B$-conserving interactions is set by the quantity $(180\\,{rm GeV}/M_f)^4 \\leq 10^{-13}$. We have shown this to be too small to allow sufficient baryogenesis for most cases. Another way to view this bound is that as the dimension-6 operator creates baryon asymmetry, the same reactants participate in creating entropy through $B$-conserving operators. The competition between the two types of reactions is almost always dominated by the $B$-conserving one. Hence, within the perturbative setting where the 4D effective action is valid, sufficient baryogenesis in this scenario is in most cases impossible almost independently of cosmology. There is one scenario which cannot simply be ruled out by the general analysis. This involves out of equilibrium $c$ quarks interacting with equilibrium $b$ quarks from a temperature of about $30\\,$GeV until a temperature of $0.2\\,$GeV. The reason why this scenario evades the suppression is that the primary entropy generation channel is through the $b$-$\\bar{b}$ self-annihilation into $c$-$\\bar{c}$ instead of the $b$-$c$ weak interaction, since $|V_{cb}|^2 \\sim 10^{-3}$. Furthermore, given that nonstandard cosmology of braneworlds may give a slow expansion rate, one can have baryogenesis persist for a long time period, giving a large integrated pileup. However, given that $c$ quarks are produced efficiently both by the gluons that keep the $b$ quarks in equilibrium as well as the $b$ quarks which undergo decay, setting up the necessary hierarchy in densities to keep $c$ out of equilibrium while $b$ in equilibrium seems difficult. However, this scenario does pass the general tests not relying upon the specifics of the out of equilibrium mechanism. Since neutrinos are almost massless and have vanishingly small tree level annihilation cross section in the limit of small $\\sqrt{s}$, one would na\\\"\\i vely think that this channel might be weak enough for the baryon number-violating operator to compete with it. Unfortunately, this turns out not to be the case. It is interesting to note that even with just the Standard Model fields, the presence of nonrenormalizable operators implies the existence of additional $CP$-violating rephasing invariants beyond the Jarlskog parameter. In any theory, the number of physical $CP$-violating phases is given by the number of couplings in the Lagrangian that are allowed to be complex, minus the dimension of the symmetry group which describes phase redefinitions of the fields: introducing new effective operators into the SM automatically creates more $CP$-violating invariants, even if all the couplings of the new operators are real in some basis. Hence, it is not the smallness of the Jarlskog invariant that prevents the success of our baryogenesis scenario. It is merely the fact that unless neutron-antineutron oscillations are suppressed by some means other than the fundamental scale, the phenomenologically acceptable scale of $100\\,$TeV is still too large for dimension-6 operators to play a significant role for baryogenesis below the temperature of the electroweak phase transition. As with any ``no go'' arguments, there are many loopholes in our conclusion (besides the one remotely possible scenario that we already mentioned). First, we neglected any possible nonperturbative physics. The 4D coupling constants are inversely proportional to the volume of the extra dimensions in which the gauge fields propagate. Thus, if the volume of the extra dimensions was small at some epoch in the early Universe, then nonperturbative physics may dominate (see Appendix \\ref{app:inst}), given that the perturbative contributions to $n_B/s$ are insufficient in our scenario. For example, the $SU(2)_L$ instanton effects may become unsuppressed, giving rise to a scenario similar to that of sphaleron transitions at the electroweak phase transition. Also, for effective temperatures below about $0.2$ GeV, non-perturbative QCD effects will be relevant, since the quark degrees of freedom become confined. Although this is not likely to change the suppression of the dimension-6 operator, the kinematics and matrix elements will be very different from our perturbative calculation done in the regime of deconfined quarks. Secondly, we have neglected all effects of the KK modes; or, in the higher dimensional picture, we have neglected the fact that the space may be inhomogeneous in the higher dimensions. For example, one may envisage a scenario in which there is a phase transition which localizes the wave functions of the quarks and the leptons. A proper description of quark-lepton separation in higher dimension will require analysis beyond the zero mode. However, this type of effect will most likely play a role in enhancing the nonequilibrium condition rather than changing the baryon number-violating branching ratio. Hence, our conclusion is most likely robust with respect to this assumption. Thirdly, we have neglected all CPT violation effects that must exist because of the time dependence of the quark-lepton separation. If the time scale associated with the quark-lepton separation process is very short, then there may be significant contributions from these effects to enhance the baryon number violating channels. However, such models will probably be severely constrained by the restriction on bulk graviton production, just as the reheating temperature is severely constrained. Hence, although there are loopholes and caveats to our ``no go'' claims, it seems fair to conclude that with only perturbative physics, degrees of freedom beyond the Standard Model fields must play a significant role in baryogenesis if the effective temperature of the Universe never exceeds 30 GeV (which is a very conservative upper bound for the maximum temperature in models with large extra dimensions, if the fundamental scale is to be accessible to collider experiments). The requirement of beyond the Standard Model field content is not an obvious statement, considering that we hardly impose any restrictions on the cosmology and that there is an unsuppressed dimension-6 operator as well as an effectively large $CP$ violating phase available for baryon asymmetry generation. Of course, our ``no go'' claims would be significantly relaxed if the dimension-9 operator responsible for neutron-antineutron oscillations is suppressed by some symmetry mechanism which is introduced by hand to supplement the geometrical suppression of proton decay operators. In closing, note that even if the fundamental scale can be lowered by evading the bound from neutron-antineutron oscillation, there is a significant challenge in this type of scenario as the quarks (and/or leptons) must be forced out of equilibrium at rather high temperatures. To accomplish this, it is possible to lower the Planck scale during baryogenesis by having a small extra dimensional volume during that period. However, there are severe restrictions coming from bulk graviton production in such cases. Furthermore, it is an extremely difficult challenge to find a time dependent potential for the scalar fields localizing the quarks and the leptons, such that the initial quark-lepton separation departs sufficiently from today's equilibrium value in a natural manner. We leave variations on our model that may generate sufficient baryon asymmetry to future studies." + }, + "0112/astro-ph0112149_arXiv.txt": { + "abstract": "We give a brief review of the physics of acoustic oscillations in Cosmic Microwave Background (CMB) anisotropies. As an example of the impact of their detection in cosmology, we show how the present data on CMB angular power spectrum on sub-degree scales can be used to constrain dark energy cosmological models. \\vspace{1pc} ", + "introduction": "\\label{intro} As it is well known, see e.g. \\cite{PZ}, Cosmic Microwave Background (CMB) anisotropies can be thought of as fluctuations $\\delta T/T$ around the mean black body temperature $T\\simeq 2.726$ K of the cosmological radiation. CMB anisotropies in a particular direction $\\hat{n}$ appear to us as a line of sight integral on the temperature fluctuations $\\delta T/T(\\hat{n},z)$ carried by CMB photons last scattered at a distance $r(z)$ from us, and weighted with the last scattering probability $P(z)$: \\begin{equation} \\label{los} {\\delta T\\over T}(here,now,\\hat{n})= \\int_{0}^{\\infty}{\\delta T\\over T}(\\hat{n},z)P(z)dz\\ . \\end{equation} The last scattering probability $P(z)$ is fixed by cosmological recombination history and it turns out to be a narrow peak around a decoupling redshift $z_{dec}$ with mean and dispersion given by $1+z_{dec}\\simeq 1100\\ ,\\ \\Delta z_{dec}\\simeq 100\\ ,$ corresponding to physical distances \\begin{equation} \\label{rdec} r_{dec}\\simeq 6000h^{-1}\\ {\\rm Mpc} \\ , \\Delta r_{dec}\\simeq 10 h^{-1}\\ {\\rm Mpc}\\ . \\end{equation} Therefore, CMB anisotropies can be thought of as a snapshot of the cosmic thermodynamical temperature in the early Universe. Their dependence on the line of sight is usually described through an expansion into spherical harmonics \\begin{equation} \\label{alm} {\\delta T\\over T}(\\hat{n}) = \\sum_{lm}a_{lm}Y_{lm}(\\hat{n})\\ . \\end{equation} The anisotropy two point correlation function, obtained averaging the product of fluctuations coming from all pairs of directions separated by an angle $\\theta$, can be expanded into Legendre polynomials $P_{l}(\\cos\\theta )$: \\begin{eqnarray} <{\\delta T\\over T}(\\hat{n}) {\\delta T\\over T}(\\hat{n'})>_{\\hat{n}\\cdot\\hat{n'}=\\cos\\theta}&=& \\nonumber\\\\ =\\sum_{l}{2l+1\\over 4\\pi}C_{l}P_{l}(\\cos\\theta )\\ &.& \\label{ctheta} \\end{eqnarray} The relation of $C_{l}$s with $a_{lm}$ coefficients is \\begin{equation} \\label{cl} C_{l}={1\\over 2l+1}\\sum_{m}|a_{lm}|^{2}\\ . \\end{equation} At high multipoles, $l\\gg 1$, the Legendre polynomials have a sharp peak at $\\theta\\simeq 200/l$ degrees; as a consequence, a $C_{l}$ coefficient quantifies the anisotropy power on the same angular scale. Moreover, taking into account that CMB anisotropies come essentially from a narrow spherical shell in redshift as in Eq.(\\ref{rdec}), also known as last scattering surface, $C_{l}$ probes perturbations on a cosmological scale $\\lambda$ represented as in figure \\ref{f1}. Given that the cosmological horizon at decoupling subtends roughly one degree on the sky, this scale separates the sub-horizon from super-horizon regimes. After the first discovery of large scale CMB anisotropies by COBE \\cite{COBE}, a breakthrough on the sub-degree structure of this signal is underway. Data from the two balloon-borne experiments BOOMERanG and MAXIMA \\cite{BOOM,MAX} and the ground based interferometer DASI \\cite{DASI} gave strong evidence of the presence of a peak at angular scales corresponding to a degree, as well as important indications for the existence of other peaks on smaller scales. Forthcoming data from satellites MAP ({\\tt http://map.nasa.gsfc.gov}) and Planck ({\\tt http://astro.estec.esa.nl/Planck}), will reveal CMB acoustic oscillations on the whole sky. In this paper we give a brief review of the most important physical mechanisms responsible for the formation of CMB acoustic peaks on sub-degree angular scales, together with some application of present data to constrain cosmological models. In Section II we put CMB anisotropies in the context of cosmological perturbation theory. In Section III we describe the phenomenology of acoustic peaks. Finally in Section IV we show an example of the impact of CMB on cosmology, briefly describing how these data can constrain cosmologies with dark energy. ", + "conclusions": "" + }, + "0112/astro-ph0112463_arXiv.txt": { + "abstract": "We decipher intrinsic three-dimensional shape distributions of molecular clouds, cloud cores, Bok globules, and condensations using recently compiled catalogues of observed axis ratios for these objects mapped in carbon monoxide, ammonia, through optical selection, or in continuum dust emission. We apply statistical techniques to compare assumed intrinsic axis ratio distributions with observed projected axis ratio distributions. Intrinsically triaxial shapes produce projected distributions which agree with observations. Molecular clouds mapped in $^{12}$CO are intrinsically triaxial but more nearly prolate than oblate, while the smaller cloud cores, Bok globules, and condensations are also intrinsically triaxial but more nearly oblate than prolate. ", + "introduction": "Numerous catalogues have now compiled properties of hundreds (or even thousands) of molecular clouds or cloud fragments with a large range of sizes, allowing meaningful statistical analysis of properties of molecular clouds, cloud cores, and smaller condensations. With such data sets, the distribution of apparent projected core axis ratio $p$ can be used to constrain the intrinsic three-dimensional shapes. In a previous paper (Jones, Basu, \\& Dubinski 2001; hereafter Paper I), we investigated the shapes of molecular cloud cores mapped in NH$_3$ (Jijina, Myers, \\& Adams 1999) and cores mapped through optical selection (Lee \\& Myers 1999) via both statistical and analytical methods. We found that strictly axisymmetric prolate or oblate shapes cores could not reproduce the observed projected axis ratios and that molecular cloud cores were triaxial. In this paper, we extend our previous work and conduct a statistical analysis of seven recent data sets which include a wide range of sizes of objects, from molecular clouds with effective radius as large as 45 pc (Heyer, Carpenter, \\& Snell 2001) to submillimeter dust continuum maps of condensations with major axes as small as 2800 AU, or $\\sim 0.01$ pc (Motte et al. 2001). The information obtained about the intrinsic shapes of these objects can yield insight into the physical processes which govern their evolution and subsequent star formation. We examine the $^{12}$CO catalogue of molecular regions in the outer Galaxy compiled by Heyer et al. (2001) in \\S\\ 3.1, cores mapped in NH$_3$ and C$^{18}$O in \\S\\ 3.2 (Onishi et al. 1996; Tachihara, Mizuno, \\& Fukui 2000), catalogues of Bok globules (Clemens \\& Barvainis 1988; Bourke et al. 1995) in \\S\\ 3.3, and millimeter and submillimeter continuum maps of smaller protostellar condensations (Motte, Andr\\'{e}, \\& Neri 1998; Motte et al. 2001) in \\S\\ 3.4. A discussion and summary are given in \\S\\S\\ 4 and 5. ", + "conclusions": "\\subsection{Error Bounds} In order to test the reliability of our results we recalculate the $\\chi^2$ values after randomly removing $20\\%$ of the data and repeat this procedure ten times for each set. For every set investigated in this paper we obtain the same best fit mean axis ratios $(\\xi_0,\\zeta_0)$ for all the trials with the exception of two trials with the Motte et al. (1998) data. This data set has only 35 non-composite objects for which the minor and major axis are published. When $20\\%$ of the data is randomly removed, only 28 objects remain. Nevertheless, we obtain $(\\xi_0,\\zeta_0) = (0.4,0.9)$ for eight trials, $(\\xi_0,\\zeta_0) = (0.5,0.9)$ for one trial, and $(\\xi_0,\\zeta_0) = (0.4,0.8)$ for the remaining trial. Additionally, adjusting the width of the Gaussian distributions in $\\xi$ and $\\zeta$ within the range $[0.05,0.2]$ yields at most a change of $\\pm 0.1$ in the best fit values $\\xi_0$ and $\\zeta_0$. The largest variation is seen in the smallest data sets. The effect on the Motte et al. (2001) data is described in \\S\\ 3.4. Using $\\sigma =0.05$ allows us to test distributions with peaks near the boundaries 0 and 1, while a relatively large $\\sigma=0.2$ allows us to see if a better fit exists for the broadest observed distributions in projected axis ratio $p$. Altogether, our testing allows us to state an approximate maximum error in our best fit mean axis ratios of $\\pm 0.1$. Interestingly, absolutely {\\em all} of the best fits for the data sets of objects with size scale $\\sim 0.1$ pc or smaller agree with one another within this range of error. This strengthens the case that these objects are triaxial but preferentially flattened in one direction and close to oblateness. A small degree of triaxiality seems necessary to explain the observed decline near $p=1$, and the pure oblate hypothesis seems excluded for dense cores (Paper I and \\S\\ 3.2). Due to the smaller number statistics and relatively large number of objects with $p$ near one, the dense condensations are most likely to be compatible with pure oblateness. Future more extensive observations of this class of objects is necessary to settle the issue. Finally, there is the possibility that our results are biased by the fact that the spectral line or dust continuum data are probing emission from regions of varying opacity and/or temperature, so that the projected shape may not correspond exactly to the physical shape of the gas distribution. When data from other wavelengths become available, it is possible our conclusions may change. This applies particularly to the molecular cloud data (\\S\\ 3.1), for which we have used a single catalogue. However, since we obtained very nearly the same result using a variety of tracers for the smaller cloud cores, Bok globules, and condensations, the results in these cases seem robust. \\subsection{Physical Implications} This investigation shows that intrinsic triaxial objects produce distributions which reasonably match observations of projected axis ratios for molecular clouds, molecular cloud cores, Bok globules, and protostellar condensations. The results clearly fall into two categories: (1) on scales $\\gtrsim 1$ pc, mapped in $^{12}$CO, molecular clouds, including GMC's, have triaxial shapes which are more closely prolate than oblate; (2) dense cores, Bok globules, and condensations, mapped in a variety of tracers, on scales from few $\\times \\,\\, 0.1$ pc down to $0.01$ pc, have triaxial shapes which are more closely oblate than prolate. The results about the latter objects reinforce our earlier finding (Paper I) from two other catalogues of dense core shapes. See Table~\\ref{summary} for a summary of the best fit axis ratios, ($\\xi_0, \\zeta_0$), for each of the data sets we investigated. The results from Paper I are included for comparison. \\begin{deluxetable}{lcrr} \\tablecolumns{4} \\tablewidth{0pc} \\tablecaption{Summary of Best Fit Mean Axis Ratios} \\tablehead{ \\colhead{Data Set} & \\colhead{Object Type} & \\colhead{$\\xi_0$} & \\colhead{$\\zeta_0$}\\\\ \\colhead{} & \\colhead{}& \\colhead{$\\pm \\, 0.1$}& \\colhead{$\\pm \\, 0.1$} } \\startdata Heyer et al. (2001)&clouds with effective radius $>$ 10 pc & 0.2& 0.2\\\\ &complete set of clouds& 0.3& 0.3\\\\ \\cline{1-4}\\\\ Onishi et al. (1996) \\& &molecular cloud cores& 0.4& 0.9\\\\ Tachihara et al. (2000)& & &\\\\ Jijina et al. (1999)\\tablenotemark{a}& molecular cloud cores & 0.5& 0.9\\\\ Lee \\& Myers (1999)\\tablenotemark{a}& molecular cloud cores & 0.3& 0.9\\\\ Clemens \\& Barvainis (1998)& Bok globules & 0.4& 0.9\\\\ Bourke et al. (1995)& Bok globules & 0.4& 0.9\\\\ Motte et al. (1998) & dense condensations & 0.4& 0.9\\\\ Motte et al. (2001) & dense condensations & 0.4& 0.9\\\\ \\enddata \\label{summary} \\tablenotetext{a}{Previous result from Paper I} \\end{deluxetable} The robust tendency for cores, Bok globules, and smaller condensations to have triaxial fits with $\\xi_0 = 0.3-0.5$, and $\\zeta_0 = 0.9$ implies that they are all preferentially flattened in {\\em one} direction. This could be due to flattening along the direction of a mean magnetic field, or due to significant rotational support in the smallest objects. The magnetic field explanation for cores is attractive as it implies that the observed near-alignment of core minor axes and magnetic field direction in Taurus (see Onishi et al. 1996) may be indicative of a more universal phenomenon. We also note that early submillimeter polarimetry of a few dense cores (Ward-Thompson et al. 2000) reveals a tendency toward alignment, but also a noticeable angular offset. This is interpreted as evidence for triaxiality of the cores (Basu 2000), which may still be preferentially flattened along the direction of the magnetic field. While triaxiality is consistent with a nonequilibrium state, evolving due to external turbulence or internal gravity, the near-oblate shape also means that the objects may not be particularly far from equilibrium, and that oblate equilibrium models may act as a reasonable approximation to these objects. This can explain why the internal structure of some Bok globules and pre-stellar cores can be closely or approximately fit by spherical equilibrium Bonnor-Ebert or near-equilibrium oblate magnetic models (Alves, Lada, \\& Lada 2001; Bacmann et al. 2000; Ciolek \\& Basu 2000; Zucconi, Walmsley, \\& Galli 2001). It is also consistent with the observed near-virial-equilibrium of most cores (Myers \\& Goodman 1988). We also note that the Bok globules, which are by definition isolated sites of star formation, have shapes that are not significantly different from that of molecular cloud cores and condensations embedded within larger clouds. This suggests that the environment in which the cores and condensations are embedded plays a relatively insignificant role in their dynamics, i.e., the external pressure from the parental cloud does not seem to be important at this stage. The smallest objects in our study, the condensations mapped in millimeter and submillimeter continuum emission, may be the precursors to individual stars since the mass spectrum appears to match the Initial Mass Function (IMF) compiled by Salpeter (1955) over a certain mass range (Motte et al. 1998). The estimated triaxial but near-oblate shape of these objects are an important link in understanding the collapse process that leads to star formation. For the larger molecular cloud scale, we have utilized an exhaustive catalogue of the shapes of clouds in the outer Galaxy (Heyer et al. 2001). Although our study of molecular cloud shapes is based on this single available sample of projected axis ratios, and the result should be confirmed when other shape data become available, the sheer size of this catalogue is a strong point. The histogram of observed axis ratios (Fig. \\ref{heyertotalbest}) is very distinct from any of the other samples. It has a very sharp peak and a severe lack of objects with $p\\gtrsim 0.5$. While there may be some unknown selection effect which biases against the observation of near-circular objects, we note that the the orientations of the projected shapes in the plane of the sky do appear to be truly random. Furthermore, we note that the earlier $^{13}$CO catalogue of only 23 clouds in Ophiuchus by Nozawa et al. (1991) that was utilized by Ryden (1996) has the same qualitative feature of a narrow peak near $p=0.3$ and a steep decline toward $p=1$. Heyer et al. (2001) note that the vast majority of their clouds (all but the largest clouds, which we loosely label GMC's) are not self-gravitating. They are either transient features or are held together by external pressure. If these clouds are indeed brought together by large scale turbulence in the interstellar medium (or even confined for some time by an anisotropic ram pressure) we might expect that they have an elongated, filamentary shape (see e.g., Nagai, Inutsuka, \\& Miyama 1998; Balsara, Ward-Thompson, \\& Crutcher 2001; review by Shu et al. 1999). Since even the largest clouds seem to have these shapes, we surmise that all clouds may be brought together by external forcing (due to shock waves or turbulent motions for example), with only the largest clouds or densest regions within smaller clouds able to become self-gravitating. This ties in with the general picture of a rapid formation of molecular clouds due to external triggers (see e.g., Hartmann, Ballesteros-Paredes, \\& Bergin 2001; Pringle, Allen, \\& Lubow 2001)." + }, + "0112/astro-ph0112180_arXiv.txt": { + "abstract": "High-resolution spectra of the hot white dwarf G191-B2B, covering the wavelength region 905--1187\\AA\\/, were obtained with the Far Ultraviolet Spectroscopic Explorer ({\\it FUSE}). This data was used in conjunction with existing high-resolution Hubble Space Telescope STIS observations to evaluate the total H~{\\sc i}, D~{\\sc i}, O~{\\sc i} and N~{\\sc i} column densities along the line of sight. Previous determinations of $N$(D~{\\sc i}) based upon GHRS and STIS observations were controversial due to the saturated strength of the D~{\\sc i} Lyman~$\\alpha$ line. In the present analysis the column density of D~{\\sc i} has been measured using only the unsaturated Lyman~$\\beta$ and Lyman~$\\gamma$ lines observed by {\\it FUSE}. A careful inspection of possible systematic uncertainties tied to the modeling of the stellar continuum or to the uncertainties in the {\\it FUSE} instrumental characteristics has been performed. The column densities derived are: $\\log N($D~{\\sc i}$)=13.40\\pm0.07$, $\\log N($O~{\\sc i}$)=14.86\\pm0.07$, and $\\log N($N~{\\sc i}$)=13.87\\pm 0.07$ quoted with $2\\sigma$ uncertainties. The measurement of the H~{\\sc i} column density by profile fitting of the Lyman~$\\alpha$ line has been found to be unsecure. If additional weak hot interstellar components are added to the three detected clouds along the line of sight, the H~{\\sc i} column density can be reduced quite significantly, even though the signal-to-noise ratio and spectral resolution at Lyman~$\\alpha$ are excellent. The new estimate of \\nhi\\ toward G191-B2B reads: $\\log N($H~{\\sc i}$)=18.18\\pm0.18$ ($2\\sigma$ uncertainty), so that the average (D/H) ratio on the line of sight is: (D/H)$=1.66^{+0.9}_{-0.6}\\times10^{-5}$ ($2\\sigma$ uncertainty). ", + "introduction": "Deuterium is thought to be produced in significant amount only during primordial Big Bang nucleosynthesis (BBN), and to be thoroughly destroyed in stellar interiors. Deuterium is thus a key element in cosmology and in Galactic chemical evolution (see e.g. Audouze \\& Tinsley 1976; Gautier \\& Owen 1983; Vidal--Madjar \\& Gry 1984; Boesgaard \\& Steigman 1985; Olive {\\it et~al.} 1990; Pagel 1992; Vangioni-Flam \\& Cass\\'e 1994; Vangioni-Flam {\\it et~al.} 1995; Prantzos 1996; Scully {\\it et~al.} 1997; Cass\\'e \\& Vangioni-Flam 1998; Tosi {\\it et al.} 1998). Indeed, its primordial abundance is the best tracer of the baryonic density parameter of the Universe, $\\Omega_B$, and the decrease in its abundance during galactic evolution traces the cosmic star formation rate at various epochs. The first (indirect) measurement of the deuterium abundance was carried out using $^3$He in the solar wind, giving the presolar value D/H$\\simeq2.5\\pm1.0\\times10^{-5}$ (Geiss \\& Reeves 1972). The first measurements of the D/H ratio in the interstellar medium (ISM) were reported shortly thereafter by Rogerson \\& York (1973), and their value (D/H)$\\simeq1.4\\pm0.2\\times10^{-5}$ has remained a landmark average value for the interstellar D/H ratio. Finally direct measurements of the primordial (D/H) ratio in low-metallicity material at high redshift have been successfully carried out these past few years (e.g., Burles 2001 for a review, and references therein). The values derived cluster around (D/H)$\\sim 3\\times10^{-5}$ although with significant dispersion, which may or may not be real. Quite similarly the measurements of the (D/H) ratio in the Galactic ISM towards hot stars with the {\\it Copernicus} satellite lead to many evaluations of D/H (see e.g. York and Rogerson 1976; Vidal--Madjar {\\it et~al.} 1977; Laurent {\\it et~al.} 1979; Ferlet {\\it et~al.} 1980; York~1983; Allen {\\it et~al.} 1992) which also show dispersion around the above York \\& Rogerson (1973) value. This dispersion has been recently confirmed by IMAPS observations (Jenkins {\\it et~al.} 1999; Sonneborn {\\it et~al.} 2000), indicating that the D/H ratio may vary by a factor $\\simeq 3$ in the solar neighborhood, i.e., within a few hundred parsecs. In this paper we present a new determination of the D/H ratio on the line of sight to the nearby DA white dwarf (WD) G191-B2B based on observations obtained with the Far Ultraviolet Spectroscopic Explorer ({\\it FUSE}, Moos {\\it et al.} 2000; Sahnow {\\it et al.} 2000). This paper is one of a series in this volume describing the first results of the {\\it FUSE} (D/H) program in the Local ISM (LISM). This program and its results are summarized in the overview paper by Moos et al. (2002). Observing white dwarfs has many advantages over hot and cool stars for studying the D/H ratio, as explained in Vidal-Madjar {\\it et al.} (1998): these targets can be chosen close to the Sun, in order to avoid a complex line of sight structure, and in the high temperature range, so that the interstellar absorption is superimposed on a smooth stellar continuum. The risk of contamination by low column density H~{\\sc i} fluffs possibly present in the hot star winds (Gry, Lamers \\& Vidal--Madjar 1984) is negligible for WDs, and their hot continuum offers the possibility of observing the numerous UV lines of N~{\\sc i} and especially O~{\\sc i}, which is a very useful tracer of H~{\\sc i} and D~{\\sc i}. The (D/H) ratio has already been measured toward four white dwarfs, using the HST: G191-B2B (Lemoine {\\it et~al.} 1996; Vidal--Madjar {\\it et~al.} 1998; Sahu {\\it et~al.} 1999), HZ43A (Landsman {\\it et~al.} 1996), Sirius~B (H\\'ebrard {\\it et~al.} 1999) and Feige~24 (Vennes {\\it et~al.} 2000). For HZ43A, Sirius~B and Feige~24 the average D/H values obtained are compatible with the local ISM (LISM) D/H determination (Linsky 1998) made within the Local Interstellar Cloud (the LIC in which the sun is embedded), although in the case of Sirius~B this compatibility is marginal. In the case of G191-B2B, it was found that the line of sight comprises one neutral region corresponding to the LIC, and two more ionized absorbing components (Lemoine {\\it et al.} 1996; Vidal-Madjar {\\it et al.} 1998). The average (D/H) ratio (defined as the ratio of the total column densities of D~{\\sc i} and H~{\\sc i}) was found to be (D/H)$=1.12\\pm0.08\\times10^{-5}$ (Vidal--Madjar {\\it et~al.} 1998), significantly different from the value measured toward Capella (D/H)$_{\\rm LIC}=1.5\\pm0.1\\times10^{-5}$ (Linsky 1998). The (D/H) ratio measured toward G191-B2B has been contested by Sahu et al. (1999), who used STIS data and concluded to the presence of two interstellar components only, and to a D/H ratio compatible with that observed toward Capella. The disagreement resides in the evaluation of the total D~{\\sc i} column density (Vidal-Madjar 2000; Sahu 2000), and arises presumably because the \\lya\\ D line is saturated and the column density is thus sensitive to the line profile. The number of components assumed on the line of sight may also introduce differences between the analyses of these groups (see Vidal-Madjar 2001 for a detailed discussion). In the present work we re-examine these issues, making use of high quality {\\it FUSE} and STIS observations of G191-B2B. We first measure the \\ndi , \\noi\\ and \\nni\\ column densities using the unsaturated lines of these elements in the {\\it FUSE} datasets, notably \\lyb\\ and \\lyg\\ for D~{\\sc i} (Section~2). We then analyze the recent high quality STIS observations in Section~3, and provide explicit evidence for the presence of three absorbing components (at least) on the line of sight. We also provide a refined estimate of the total H~{\\sc i} column density. All throughout this work, considerable effort has been put on quantifying possible systematic uncertainties related to fixed-pattern noise, detector artifacts, background uncertainties, wavelength calibration and modeling of the stellar continua, as well as to the velocity structure of the line of sight. In particular, we argue in Section~3.3 that previous estimates of the total \\nhi\\ are subject to a large systematic uncertainty related to the possible presence of additional weak [\\nhi$\\leq10^{14}\\,$cm$^{-2}$] hot ($T\\sim10^5\\,$K) components. This effect may have a large impact on our understanding of the observed variations of the (D/H) ratio in the ISM, as it may affect other lines of sight, and is the subject of a companion paper (Vidal-Madjar \\& Ferlet 2002). We provide a summary of our results and a short discussion in Section~4; an overall discussion of the {\\it FUSE} results is given by Moos {\\it et al.} (2002). ", + "conclusions": "We have measured the total column densities of D~{\\sc i}, N~{\\sc i} and O~{\\sc i} toward G191-B2B using unsaturated absorption lines of these elements in high quality {\\it FUSE} spectra. After a careful examination of the possible systematic uncertainties tied to the choice of the stellar continuum and to the instrumental configuration, we have derived the following column densities with $2\\sigma$ uncertainties: \\begin{center} $\\log N({\\rm D}$~{\\sc i}$)_{\\rm tot} = 13.40 \\pm0.07,$ \\\\ $\\log N({\\rm O}$~{\\sc i}$)_{\\rm tot} = 14.86 \\pm0.07,$ \\\\ $\\log N({\\rm N}$~{\\sc i}$)_{\\rm tot} = 13.87 \\pm0.07.$ \\end{center} We have also analyzed new high signal-to-noise ratio high resolution STIS observations of G191-B2B and provided concrete evidence for the presence of at least three interstellar absorbing components on the line of sight by analyzing the interstellar absorption lines of N~{\\sc i}, O~{\\sc i}, Si~{\\sc ii}, Si~{\\sc iii}, S~{\\sc iii} and Fe~{\\sc ii} present in the STIS bandpass. We have also measured the total hydrogen column density on the line of sight using the velocity structure derived from the above metals. We have performed an exhaustive study of systematic effects on the value of \\nhi. In particular we have discovered a new major source of uncertainty on \\nhi\\ tied to the possible presence of additional weak hot absorbers whose combined absorption profile can contribute significantly to the wings of the blended \\lya\\ profile. The column density of these absorbers is small compared to the other main components, and they would not be detected in any other species than H~{\\sc i}, but their contribution to the Lyman~$\\alpha$ absorption profile can reduce significantly the total H~{\\sc i} column density measured from the profile fitting. In order to constrain their impact, we have analyzed simultaneously Lyman~$\\alpha$ and the higher order Lyman lines, and concluded that the best value of \\nhi\\ toward G191-B2B is: \\begin{center} $\\log N({\\rm H}$~{\\sc i}$)_{\\rm tot} = 18.18 \\pm 0.18 \\qquad (2\\sigma\\,{\\rm error})$ \\end{center} We emphasize that this uncertainty is a systematic uncertainty which had gone unnoticed before. Therefore the above result supersedes previous estimates of \\nhi\\ toward G191-B2B obtained from the profile fitting of \\lya. A detailed analysis of this uncertainty and its consequences on \\nhi\\ determinations toward other stars is discussed in a companion paper (Vidal-Madjar \\& Ferlet 2002). We thus derive the following neutral abundance ratios toward G191-B2B, with $2\\sigma$ uncertainties: \\begin{eqnarray} ({\\rm D}/{\\rm H})_{\\rm tot}\\, & = &\\, 1.66^{+0.9}_{-0.6}\\, 10^{-5}\\nonumber \\\\ ({\\rm D}/{\\rm O})_{\\rm tot}\\, & = &\\, 3.49 \\pm 0.78 \\, 10^{-2}\\nonumber \\\\ ({\\rm D}/{\\rm N})_{\\rm tot}\\, & = &\\, 3.41 \\pm 0.76 \\, 10^{-1}\\nonumber \\\\ ({\\rm O}/{\\rm H})_{\\rm tot}\\, & = &\\, 4.79^{+2.5}_{-1.7}\\, 10^{-4}\\nonumber \\\\ ({\\rm N}/{\\rm H})_{\\rm tot}\\, & = &\\, 4.90^{+2.6}_{-1.8}\\, 10^{-5}\\nonumber \\\\ \\end{eqnarray} Most of the uncertainty in the above result results from the systematic uncertainty on the \\nhi\\ determination. This clearly shows the importance of measuring accurate (D/O) and (D/N) ratios in the interstellar medium instead of abundances relative to hydrogen, as emphasized by Timmes {\\it et al.} (1997). Interestingly if one uses the recent measurement of \\nhi\\ from the modeling of the atmosphere of G191-B2B and the fit of the EUVE spectrum, $\\log N$(H~{\\sc i})$=18.30\\pm0.09$ ($2\\sigma$), one finds $({\\rm D}/{\\rm H})_{\\rm tot}=1.26_{-0.29}^{+0.36}\\times 10^{-5}$ ($2\\sigma$). At this stage, however, due to the uncertainty inherent to the modeling of the white dwarf atmosphere, it is probably more conservative to use the interstellar determination for \\nhi, and therefore the previous value of the (D/H) ratio. The above new value for the (D/H) ratio agree with the range of values measured by Linsky (1998) toward a dozen stars of the LISM and with the values previously derived toward G191-B2B. However the discrepancy between previous estimates of the (D/H) ratio toward G191-B2B and the LISM average D/H ratio has disappeared due to a revision of the uncertainty on the estimation of the total H~{\\sc i} content. A detailed interpretation of this (D/H) value and of the accompanying (D/O) and (D/N) ratios and their implications is provided in a companion paper by Moos {\\it et al.} (2002)." + }, + "0112/astro-ph0112239_arXiv.txt": { + "abstract": "After a couple of years of quiescence, the soft gamma repeater SGR 1900+14 suddenly reactivated on 18 April 2001, with the emission of a very intense, long and modulated flare, only second in intensity and duration to the 27 August 1998 giant flare. BeppoSAX caught the large flare with its Gamma Ray Burst Monitor and with one of the Wide Field Cameras. The Wide Field Cameras also detected X-ray bursting activity shortly before the giant flare. A target of opportunity observation was started only 8 hours after the large flare with the Narrow Field Instruments, composed of two 60-ks long pointings. These two observations show an X-ray afterglow of the persistent SGR 1900+14 source, decaying with time according to a power law of index -0.6. ", + "introduction": "On April 18, 2001 at 07:55:11.509 UT the Gamma Ray Burst Monitor (GRBM, 40-700 keV) onboard BeppoSAX was triggered by a large flare. The event was also observed by the unit 1 of the BeppoSAX Wide Field Cameras (WFCs, 2-28 keV), that was automatically shut-down few seconds after the start of the event by an onboard safeguard algorithm. The WFC localized this giant flare as originating from the soft gamma-ray repeater SGR 1900+14. The BeppoSAX GRBM unit 1, co-aligned to WFC unit 1, detected a peak flux of 16,400 counts/s (40-700 keV, after deadtime correction), corresponding to (1.1$\\pm$0.1)$\\times10^{-5}$ erg cm$^{-2}$ s$^{-1}$. The event lasted approximately 40~s, and the total fluence was about 240,000 counts. A 5-s periodic modulation of the light curve is consistent with the 5.17-s period of the persistent X-ray source. In the light curve of the flare we can identify the repetition of 5 or 6 cycles, and a last pulse that is out of phase with respect to the previous ones. Assuming a thermal bremsstrahlung spectral shape, from our 2-channel, 1-s resolved spectrum, we can derive a plasma temperature kT of $\\sim$30~keV. There is no evidence for a hard spike at the start of the event, and the spectral evolution during the event is relatively smooth, with a general hardening trend. The Konus-Wind instrument detected this event in the non-trigger mode because of very high background level after strong solar flare (30 times more than normal). Our estimation for energy more than 15 keV is: fluence - 2$\\times$10$^{-4}$ erg cm$^{-2}$ and peak flux - 2.5$\\times$10$^{-5}$ erg cm$^{-2}$ s$^{-1}$. In spite of rough time resolution in the background mode the Konus-Wind data has evidences of the 5-s pulsations. Three weak and short bursts are detected at times 2537, 755 and 444 seconds {\\it before} the giant flare in the WFC unit 1. Their durations are 100, 125 and 55 ms, and their peak fluxes exceed 20 Crab units in 2-28 keV (Feroci et al. 2001b). No gamma-ray counterpart was found for these events in the GRBM, nor in {\\it Ulysses} or Konus/Wind. This is the first reported bursting activity from SGR 1900+14 since approximately two years. At the time of the giant flare, as detected by the GRBM, the WFC1 detected a large increase in the count rate for approximately 3 seconds after the GRBM trigger. At this time the count rate was so high (larger than 30,000 counts/s) that it triggered the automatic turn off of the experiment. ", + "conclusions": "The BeppoSAX instrumentation observed the most interesting phases of the 2001 activation of SGR 1900+14 with a very complete set of observations. The GRBM and WFC caught the second giant flare in the (observed) history of this source and its precursor X-ray activity. The NFI observed the afterglow of the persistent source after the flare, the bursting activity and its quenching, and the period history soon after the giant flare. The giant flare is rather unusual with respect to its two predecessors, the March 5 1979 event from SGR 0526-66 and the August 27 1998 flare from SGR 1900+14 (e.g., Mazets et al. 1979, Hurley et al. 1999, Feroci et al. 2001a). The most important differences in the April 18 are in the following aspects: \\begin{itemize} \\item the peak flux and fluence are more than one order of magnitude smaller \\item the duration is almost one order of magnitude shorter \\item there is no evidence for the first short and very hard pulse at the beginning of the event, implying that, contrary to the other two events, most likely there was no relativistic particle emission during this event. \\end{itemize} The NFI observation is the fastest pointing after a giant flare and provides a {\\it direct} observation of a long-duration afterglow emission from an SGR after a giant flare. Our result of an afterglow showing a power law decay, with index -0.6, is similar to that obtained by Woods et al. (2001), who measured a decay index of -0.713$\\pm$0.025 after the August 27 event, based on the assumed constancy of the pulsed fraction. Interestingly, Ibrahim et al. (2001) found a similar decay law in the pulsating soft tail of the August 29 short burst. The spectral variability of the X-ray afterglow is also interesting. During our first observation, at high flux level, the source spectrum showed a softening, and an overall power law shape. In our second observation, at a flux typical of the quiescent state, we clearly detect the emergence of a kT=0.5 keV blackbody component, and an N$_{H}$ smaller by about a factor of 2. Based on our observations, together with our analysis of the previous BeppoSAX observations (see also Woods et al. 1999, 2001) of this source in quiescence (May 1997 and March/April 2000) and after the August 27 giant flare (September 1998), we can state that the blackbody component for this source is peculiar of its quiescent state." + }, + "0112/astro-ph0112525_arXiv.txt": { + "abstract": "{A number of studies of abundance gradients in the galactic disk have been performed in recent years. The results obtained are rather disparate: from no detectable gradient to a rather significant slope of about $-0.1$ dex~kpc$^{-1}$. The present study concerns the abundance gradient based on the spectroscopic analysis of a sample of classical Cepheids. These stars enable one to obtain reliable abundances of a variety of chemical elements. Additionally, they have well determined distances which allow an accurate determination of abundance distributions in the galactic disc. Using 236 high resolution spectra of 77 galactic Cepheids, the radial elemental distribution in the galactic disc between galactocentric distances in the range 6-11 kpc has been investigated. Gradients for 25 chemical elements (from carbon to gadolinium) are derived. The following results were obtained in this study: Almost all investigated elements show rather flat abundance distributions in the middle part of galactic disc. Typical values for iron-group elements lie within an interval from $\\approx -0.02$ to $\\approx -0.04$ dex~kpc$^{-1}$ (in particular, for iron we obtained d[Fe/H]/dR$_{\\rm G}= -0.029$ dex~kpc$^{-1}$). Similar gradients were also obtained for O, Mg, Al, Si, and Ca. For sulphur we have found a steeper gradient ($-0.05$ dex~kpc$^{-1}$). For elements from Zr to Gd we obtained (within the error bars) a near to zero gradient value. This result is reported for the first time. Those elements whose abundance is not expected to be altered during the early stellar evolution (e.g. the iron-group elements) show at the solar galactocentric distance [El/H] values which are essentially solar. Therefore, there is no apparent reason to consider our Sun as a metal-rich star. The gradient values obtained in the present study indicate that the radial abundance distribution within $\\approx$ 10 kpc is quite homogeneous, and this result favors a galactic model including a bar structure which may induce radial flows in the disc, and thus may be responsible for abundance homogenization. ", + "introduction": "In recent years the problem of radial abundance gradients in spiral galaxies has emerged as a central problem in the field of galactic chemodynamics. Abundance gradients as observational characteristics of the galactic disc are among the most important input parameters in any theory of galactic chemical evolution. Further development of theories of galactic chemodynamics is dramatically hampered by the scarcity of observational data, their large uncertainties and, in some cases, apparent contradictions between independent observational results. Many questions concerning the present-day abundance distribution in the galactic disc, its spatial properties, and evolution with time, still have to be answered. Discussions of the galactic abundance gradient, as determined from several studies, were provided by Friel (\\cite{fr95}), Gummersbach et al. (\\cite{guet98}), Hou, Prantzos \\& Boissier (\\cite{hpb00}). Here we only briefly summarize some of the more pertinent results. 1) A variety of objects (planetary nebulae, cool giants/supergiants, F-G dwarfs, old open clusters) seem to give evidence that an abundance gradient exists. Using DDO, Washington, UBV photometry and moderate resolution spectroscopy combined with metallicity calibrations for open clusters and cool giants the following gradients were derived (d[Fe/H]/dR$_{\\rm G}$): $-0.05$ dex~kpc$^{-1}$ (Janes \\cite{jan79}), $-0.095$ dex~kpc$^{-1}$ (Panagia \\& Tosi \\cite{pt81}), $-0.07$ dex~kpc$^{-1}$ (Harris \\cite{har81}), $-0.11$ dex~kpc$^{-1}$ (Cameron \\cite{cam85}), $-0.017$ dex~kpc$^{-1}$ (Neese \\& Yoss \\cite{ny88}), $-0.13$ dex~kpc$^{-1}$ (Geisler, Clari\\'a \\& Minniti \\cite{gcm92}), $-0.097$ dex~kpc$^{-1}$ (Thogersen, Friel \\& Fallon \\cite{tff93}), $-0.09$ dex~kpc$^{-1}$ (Friel \\& Janes \\cite{fj93}), $-0.091$ dex~kpc$^{-1}$ (Friel \\cite{fr95}), $-0.09$ dex~kpc$^{-1}$ (Carraro, Ng \\& Portinary \\cite{cnp98}), $-0.06$ dex~kpc$^{-1}$ (Friel \\cite{fr99}, Phelps \\cite{ph00}). One must also add that there have been attempts to derive the abundance gradient (specifically d[Fe/H]/dR$_{\\rm G}$) using high-resolution spectroscopy of cool giant and supergiant stars. Harris \\& Pilachowski (\\cite{hp84}) obtained $-0.07$ dex~kpc$^{-1}$, while Luck (\\cite{luck82}) found a steeper gradient of $-0.13$ dex~kpc$^{-1}$. Oxygen and sulphur gradients determined from observations of planetary nebulae are $-0.058$ dex~kpc$^{-1}$ and $-0.077$ dex~kpc$^{-1}$ respectively (Maciel \\& Quireza \\cite{mq99}), with slightly flatter values for neon and argon, as in Maciel \\& K\\\"oppen (\\cite{mk94}). A smaller slope was found in an earlier study of Pasquali \\& Perinotto (\\cite{pp93}). According to those authors the nitrogen abundance gradient is $-0.052$ dex~kpc$^{-1}$, while that of oxygen is $-0.030$ dex~kpc$^{-1}$. 2) From young B main sequence stars, Smartt \\& Rolleston (\\cite{sr97}) found a gradient of $-0.07$ dex~kpc$^{-1}$, while Gehren et al. (\\cite{gehet85}), Fitzsimmons, Dufton \\& Rolleston (\\cite{fdr92}), Kaufer et al. (\\cite{kaet94}) and Kilian-Montenbruck, Gehren \\& Nissen (\\cite{kmgn94}) derived significantly smaller values: $-0.03-0.00$ dex~kpc$^{-1}$. No systematic abundance variation with galactocentric distance was found by Fitzsimmons et al. (\\cite{fitet90}). The recent studies of Gummersbach et al. (\\cite{guet98}) and Rolleston et al. (\\cite{rollet00}) support the existence of a gradient ($-0.07$ dex~kpc$^{-1}$). The elements in these studies were C-N-O and Mg-Al-Si. 3) Studies of the abundance gradient (primarily nitrogen, oxygen, sulphur) in the Galactic disc based on young objects such as \\ion{H}{ii} regions give positive results: either significant slopes from $-0.07$ to $-0.11$ dex~kpc$^{-1}$ according to: Shaver et al. (\\cite{shavet83}) for nitrogen and oxygen, Simpson et al. (\\cite{simpet95}) for nitrogen and sulphur, Afflerbach, Churchwell \\& Werner (\\cite{afflet97}) for nitrogen, Rudolph et al. (\\cite{rudet97}) for nitrogen and sulphur, or intermediate gradients of about $-0.05$ to $-0.06$ dex~kpc$^{-1}$ according to: Simpson \\& Rubin (\\cite{simru90}) for sulphur, Afflerbach, Churchwell \\& Werner (\\cite{afflet97}) for oxygen and sulphur; and negative ones: weak or nonexistent gradients as concluded by Fich \\& Silkey \\cite{fs91}; Vilchez \\& Esteban \\cite{viles96}, Rodriguez (\\cite{rod99}). Recently Pe\\~na et al. (\\cite{pet00}) derived oxygen abundances in several \\ion{H}{ii} regions and found a rather flat distribution with galactocentric distance (coefficient $-0.04$ dex~kpc$^{-1}$). The same results were also reported by Deharveng et al. (\\cite{deet00}). As one can see, there is no conclusive argument allowing one to come to a definite conclusion about whether or not a significant abundance gradient exists in the galactic disc, at least for all elements considered and within the whole observed interval of galactocentric distances. Compared to other objects supplying us with an information about the radial distribution of elemental abundances in the galactic disc, Cepheids have several advantages: 1) they are primary distance calibrators which provide excellent distance estimates; 2) they are luminous stars allowing one to probe to large distances; 3) the abundances of many chemical elements can be measured from Cepheid spectra (many more than from \\ion{H}{ii} regions or B stars). This is important for investigation of the distribution in the galactic disc of absolute abundances and abundance ratios. Additionally, Cepheids allow the study of abundances past the iron-peak which are not generally available in \\ion{H}{ii} regions or B stars; 4) lines in Cepheid spectra are sharp and well-defined which enables one to derive elemental abundances with high reliability. In view of the inconsistencies in the current results on the galactic abundance gradient, and those advantages which are afforded by Cepheids, we have undertaken a large survey of Cepheids in order to provide independent information which should be useful as boundary conditions for theories of galactic chemodynamics. We also hope that the results on the abundance gradient from the Cepheids will also be helpful to constrain the structure and age of the bar, and its influence on the metallicity gradient. This first paper in this series on abundance gradients from Cepheids presents the results for the solar neighbourhood. ", + "conclusions": "\\subsection{The radial distribution of elemental abundances: general picture and remarks on some elements} Using our calculated galactocentric distances and average abundances we can determine the galactic metallicity gradient from a number of species. Plots for several chemical elements and results of a linear fit are given in Fig. 5 (iron) and Figs. 6-9 (other elements). Note that, in the plots for Si and Cr, TX Del is not included. This star shows rather strong excess in the abundances of these elements which could be connected with its peculiar nature (in Harris \\& Welch \\cite{hw89} TX Del is reported as a spectroscopic binary. It has also been labeled a Type II Cepheid at times). In the plot for carbon we did not include the data for FN Aql and SV Mon, both of which have an extremely low carbon abundances. These two unusual Cepheids will be discussed in detail in a separate paper. The information in the plots and also in Fig. 10 enables one to put together several important conclusions. Most radial distributions of the elements studied indicate a negative gradient ranging from about $-0.02$ dex~kpc$^{-1}$ to $-0.06$ dex~kpc$^{-1}$, with an average of $-0.03$ dex~kpc$^{-1}$ for the elements in Figs. 5-8. The most reliable value comes from iron (typically the number of iron lines for each star is about 200-300). The gradient in iron is $-0.029$ dex~kpc$^{-1}$, which is close to the typical gradient value produced by other iron-group elements. Examination of Fig. 5 might lead one to suspect that the iron gradient is being controlled by the cluster of stars at R $\\approx 6.5$ with [Fe/H] $\\approx 0.2$. If one deletes these stars from the solution the gradient falls to approximately $-0.02$ dex~kpc$^{-1}$. This latter value differs from the value determined using all the data by about twice the formal uncertainty in either slope. However, we do not favour the neglect these points as there is no reason to suspect these abundances relative to the bulk of the objects. Indeed, in a subsequent paper, we shall present results for Cepheids which lie closer to the galactic center and which have abundances above those of this study, which may imply a steepening of the gradient towards the galactic center. Unweighted iron abundances give a gradient of $-0.028$ dex~kpc$^{-1}$. Both weighted and unweighted iron gradients are not significantly changed if we remove two Cepheids at galactocentric distances greater than 11 kpc (gradient is $-0.031$ dex~kpc$^{-1}$). Thus, the average slope of about $-0.03$ dex~kpc$^{-1}$ probably applies to the range $6 \\le$ R$_{\\rm G}$ (kpc) $\\le 10$. Notice that in all cases the correlation coefficient is relatively low, r $\\approx 0.47$. Carbon shows a surprisingly clear dependence upon galactocentric distance (Fig. 6a): the slope of the relation is among the largest from examined elements. We have included in the present study elements such as carbon and sodium, although the gradients based on their abundances determined from Cepheids may not be conclusive. In fact, it is quite likely that the surface abundances of these elements have been altered in these intermediate mass stars during their evolution from the main sequence to the Cepheid stage. For example, the surface abundance of carbon should be decreased after the global mixing which brings the CNO-processed material into the stellar atmosphere (turbulent diffusion in the progenitor B main sequence star, or the first dredge-up in the red giant phase). Some decrease in the surface abundance of oxygen is also expected for supergiant stars, but at a significantly lower level than for carbon (Schaller et al. \\cite{schallet92}). It is also well known that galactic supergiants (Cepheids, in particular) show an increased sodium abundance which is usually interpreted as a result of dredge-up of material processed in the Ne--Na cycle (and therefore enriched in sodium) to the stellar surface (Sasselov \\cite{sass86}, Luck \\cite{luck94}, Denissenkov \\cite{denis93a}, \\cite{denis93b}, \\cite{denis94}). Such a contamination of the Cepheids' atmospheres with additional sodium may result in a bias of the [Na/H] gradient value derived from Cepheids in relation to the true gradient. It can be seen that our results in Fig. 6a,c are consistent with these considerations on C and Na, respectively. It is not clear how these effects should affect stars at different galactocentric distances (with different metallicities), but it is likely that they contribute to increase the dispersion in the abundances, thus producing a flatter gradient. There are some indications (Andrievsky \\& Kovtyukh \\cite{ankov96}) that surface Mg and Al abundances in yellow supergiants can be altered to some extent due to mixing of the material processed in the Mg--Al cycle with atmospheric gas. This supposition seems to gain some additional support from our present data (see Fig.11) where one can see that the Mg and Al abundances are correlated. As surface abundance modifications depend upon the number of visits to the red giant region (i.e. the number of dredge-up events) as well as other factors (pre dredge-up events, depth of mixing events, mass), it is possible that the program Cepheids could show differential evolutionary effects in their abundances. Because of the high probability of such effects impacting the observed carbon and sodium (and perhaps, oxygen, magnesium, and aluminum) abundances in these Cepheids, we recommend that our gradient values for carbon and sodium to be viewed with extreme caution, while the gradients of oxygen, magnesium and aluminum abundances could be used, but also with some caution. The difference in metallicity between the stars of our sample (say, at 6 kpc and 10 kpc) is about 0.25 dex. This is a rather small value to detect/investigate the so-called \"odd-even\" effect, that is the metallicity dependent yield for some elements which should be imprinted on the trends of abundance ratios for [El$_{\\rm odd}$/El$_{\\rm even}$] versus galactocentric distance, see for details Hou, Prantzos \\& Boissier (\\cite{hpb00}). Such elements as, for example, aluminum, scandium, vanadium and manganese should show progressively decreasing abundances with overall metal decrease. This should manifest itself as a gradient in [El$_{\\rm odd}$/Fe]. We have plotted the abundance ratios for some \"odd\" elements (normalized to iron abundance) versus R$_{\\rm G}$ in Fig.12. As one can see, none of the abundance ratios plotted versus galactocentric distance shows a clear dependence upon R$_{\\rm G}$. This could mean that the \"odd-even\" effect may be overestimated if only the yields from massive stars are taken into account ignoring other possible sources, or that the effect is not sufficiently large to be seen over the current distance and metallicity baseline. \\subsection{Metallicity dispersion and the metallicity in the solar vicinity} There is a spread in the metallicity at each given galactocentric distance (larger than the standard error of the abundance analysis) which is most likely connected with local inhomogeneities in the galactic disc. As an example, in Fig. 13, we show the derived iron abundance vs. galactic longitude for the stars of our sample (a few Cepheids with heliocentric distances large than 3000 pc were excluded). The distribution gives only a small hint about a local increase of the metallicity in the solar vicinity towards the direction l$\\approx 30^{\\degr}$ and $150^{\\degr}$. It is important to note that at the solar galactocentric distance those elements, whose abundance is not supposed to be changed in supergiants during their evolution, show on average the solar abundance in Cepheids. Relative to the solar region, the stars within our sample which are within 500 pc of the Sun have a mean [Fe/H] of $\\approx +0.01$ (n = 14, $\\sigma = 0.06$). If we consider all program stars at a galactocentric radius of 7.4--8.4 kpc, i.e. those in a 1 kpc wide annulus centered at the solar radius, we find a mean [Fe/H] of approximately +0.03 (n = 29, $\\sigma = 0.05$). This result again stresses the importance of the problem connected with subsolar metallicities reported for the hot stars from the solar vicinity (see, e.g. Gies \\& Lambert \\cite{gila92}, Cunha \\& Lambert \\cite{cunla94}, Kilian \\cite{kil92}, Kilian, Montenbruck \\& Nissen \\cite{kmn94}, Daflon, Cunha \\& Becker \\cite{dafcube99}, Andrievsky et al. \\cite{anet99}). This also follows from the plots provided by Gummersbach et al. (\\cite{guet98}) for several elements. This problem was discussed, for instance, by Luck et al. (\\cite{lucket00}). The authors compared the elemental abundances of B stars from the open cluster M25 with those of the Cepheid U Sgr and two cool supergiants which are also members of the cluster, and found disagreement in the abundances of the B stars and supergiants; e.g., while the supergiants of M25 show nearly solar abundances, the sample of B stars demonstrate a variety of patterns from under- to over-abundances. This should not be observed if we assume that all stars in the cluster were born from the same parental nebula. Obviously, the problem of some disagreement between abundance results from young supergiants and main-sequence stars requires further investigation. \\subsection{Flattening of the elemental distribution in the solar neighbourhood} All previous studies of the radial abundance distribution in the galactic disc have considered only chemical elements from carbon to iron, and all derived gradients have shown a progressive decrease in abundance with increasing galactocentric distance. For the elements from carbon to yttrium in this study our gradient values also have negative signs, while for the heavier species (from zirconium to gadolinium) we obtained (within the error bars) near-to-zero gradients (see Fig. 10). Two obvious features which are inherent to derived C-Gd abundance distributions have to be interpreted: a rather flat character of the distribution for light/iron-group elements, and an apparent absence of a clear gradient for heavy species. The flattening of the abundance distribution can be caused by radial flows in the disc which may lead to a homogenization of ISM. Among the possible sources forcing gas of ISM to flow in the radial direction, and therefore producing a net mixing effect there could be a gas viscosity in the disc, gas infall from the halo, gravitational interaction between gas and spiral waves or a central bar (see e.g., Lacey \\& Fall \\cite{lf85} and Portinari \\& Chiosi \\cite{pch00}). The mechanism of the angular momentum re-distribution in the disc based on the gas infall from the halo is dependent upon the infall rate, and therefore it should have been important at the earlier stages of the Galaxy evolution, while other sources of the radial flows should effectively operate at present. Gravitational interaction between the gas and density waves produces the radial flows with velocity (Lacey \\& Fall \\cite{lf85}): \\begin{equation} v_{\\rm r} \\sim (\\Omega_{\\rm p} - \\Omega_{\\rm c})^{-1}, \\end{equation} where $\\Omega_{\\rm p}$ and $\\Omega_{\\rm c}$ are the angular velocities of the spiral wave and the disc rotation respectively. According to Amaral \\& L\\'epine (\\cite{amle97}) and Mishurov et al. (\\cite{miet97}) among others, based on several different arguments, the galactic co-rotation resonance is located close to the solar galactic orbit. The co-rotation radius is the radius at which the galactic rotation velocity coincides with the rotation velocity of the spiral pattern. Together with Eq. (4) this means that, inside the co-rotation circle, gas flows towards the galactic center ($\\Omega_{\\rm c} > \\Omega_{\\rm p}$ and v$_{\\rm r} < 0$), while outside it flows outwards. This mechanism can produce some \"cleaning\" effect in the solar vicinity, and thus can lead to some flattening of the abundance distribution. In addition, it could explain the similarity in the solar abundances and mean abundances in the five billion years younger Cepheids located at the solar galactocentric annulus (see Fig. 5), although one might expect that the Cepheids from this region should be more abundant in metals than our Sun. There is a clear evidence that the bars of spiral galaxies have also a great impact on chemical homogenization in the discs (Edmunds \\& Roy \\cite{er93}; Martin \\& Roy \\cite{mr94}, Gadotti \\& Dos Anjos \\cite{gda01}). It has been shown that a flatter abundance gradient is inherent to galaxies which have a bar structure. This could imply that a rotating bar is capable of producing significant homogenization of the interstellar medium, while such homogenization is not efficient in unbarred spiral galaxies. The direct detection of a bar at the center of our galaxy using $COBE$ maps was reported by Blitz \\& Spergel (\\cite{bs91}). Kuijken (\\cite{kui96}), Gerhard (\\cite{ger96}), Gerhard, Binney \\& Zhao (\\cite{geret98}), Raboud et al. (\\cite{rabet98}) also suggest that the Milky Way is a barred galaxy. The most recent evidence for a long thin galactic bar was reported by L\\'opez-Corredoira et al. (\\cite{lcet01}) from the DENIS survey. These authors conclude that our Galaxy is a typical barred spiral. If so, then the Milky Way should obey the relation between the slope of metallicity distribution and the bar strength (specifically, the axial ratio), which is based on the data obtained from other galaxies. According the above mentioned authors the galactic bar is triaxial and has an axial ratio (b/a) of about $1/3-1/2$ (see also Fux \\cite{fux97}, \\cite{fux99}). With such axial ratio an ellipticity E$_{\\rm B}$ = 10 (1-b/a) $\\approx 5-7$. L\\'epine \\& Leroy (\\cite{ll00}) presented a model which reproduces a near-infrared brightness distribution in the Galaxy. Their estimate of the galactic bar characteristics supposes that the total length of the bar should be about 4.6 kpc, while its width about 0.5 kpc. In this case an ellipticity could be even larger than 7. For such ellipticities the observational calibration of Martin \\& Roy (\\cite{mr94}) for barred galaxies predicts a metallicity slope of about $-0.03$ to $\\approx 0$ dex~kpc$^{-1}$ for oxygen. Our results on abundance gradients in the solar neighbourhood for iron-group elements and light species (such as Si, Ca, and even oxygen) appear to be in good agreement with expected gradient value which is estimated for the galactic disc solely from bar characteristics. Martinet \\& Friedli (\\cite{mf97}) investigated secular chemical evolution in barred systems and found that a strong bar is capable of producing significant flattening of the initial gradient across the disc. Using numerical results of that paper one can trace the (O/H) abundance evolution in barred systems. With our abundance gradients for such elements as oxygen, silicon, calcium and iron-group elements one can conclude that an expected age of the galactic bar is approximately 1 Gyr, or less. Another important result obtained by Martinet \\& Friedli (\\cite{mf97}) is that the bar of such an age should produce not only significant flattening across almost the whole disc, but also steepening of the abundance distribution in the inner parts (our observational results for this region will be discussed in the next paper from this series). An additional mechanism which may cause some local flattening (or even a shallow local minimum in the elemental abundances) should operate near the galactocentric solar radius where the relative rotational velocity of the disc and spiral pattern is small. The shocks that arise when the gas orbiting in the disc penetrates the spiral potential perturbation, and which are responsible for triggering star formation in spiral arms, pass through a minimum strength at this galactic radius, due to almost zero relative velocity. Furthermore, simulations performed by L\\'epine, Mishurov \\& Dedikov (\\cite{lmd01}) show that there is also a gas depletion at the co-rotation radius. Both reasons point towards a minimum of star formation rate at the co-rotation radius. This lower star formation rate manifests itself in the models as a minimum in elemental abundances. One can expect that after a few billion years, a galactic radius with minimum star formation rate should correspond to a local minimum in metallicity. The flat local minimum in metal abundance should be observable, unless the mechanisms that produce radial transport or radial mixing of the gas in the disc are important, or if the co-rotation radius varied appreciably in a few billion years. Note that the star-formation rate also depends on the gas density, which decreases towards large galactic radii. The combined effect of gas density and co-rotation could produce a slightly displaced minimum. At first glance, the abundance data presented in Figs. 5-9 show little indication of a local abundance minimum (or discontinuity) at the solar galactocentric radius. Nevertheless, the parabolic fit of the iron abundance distribution rather well represents observed data, and shows that a small increase in the metallicity at galactic radii larger than the co-rotation radius may not be excluded (Fig. 14). Comparing gradients from iron-group elements (small and negative) with those from the heaviest species (near to zero) one could propose the following preliminary explanation of the observed difference. The known contributors of the O-to-Fe-peak nuclei to ISM are massive stars exploding as SNe II (short-lived) and SNe I (long-lived), while s-process elements (past iron-peak) are created only in the low-mass AGB stars (1-4 M$_{\\odot}$, Travaglio et al. \\cite{travet99}). The extremely flat distribution in the disc seen for s-process elements implies that there should exist some mechanism(s) effectively mixing ISM at time-scales less than the life times of the stars with masses 1-4 M$_{\\odot}$ ($\\tau \\approx 0.3-10$ Gyr). At the same time such a mechanism may not be able to completely erase the O-Fe gradients related to the ISM, and imprinted on the young stars. If the characteristic time of the mixing (even being possibly comparable to the SNe I life time) exceeds a nuclear evolution of the SNe II O-Fe contributors, then these are the high-mass stars that could be responsible for the resulting small negative gradients from O-Fe elements in the disc. If one adopts the velocity of the radial flows, say, 4 km~s$^{-1}$ (see discussion in Lacey \\& Fall \\cite{lf85} and Stark \\& Brand \\cite{stbr89}), then the necessary time to mix the gas within about 4 kpc (baseline covered by our data) should be likely less than 1 Gyr, that is below the life-time intervals for AGB progenitors with 1-2 M$_{\\odot}$. However, this $ad~hoc$ supposition meets a problem with the observed Eu gradient. This element is believed to be produced mainly through the r-process in lower-mass SNe II (e.g., Travaglio et al. \\cite{travet99}), and therefore should probably behave similar to, for example, iron, but its radial abundance distribution appears to be quite similar to that of the s-process elements, like Zr, La, Ce, Nd (see Fig. 10)." + }, + "0112/astro-ph0112531_arXiv.txt": { + "abstract": "We study the local behavior of gravitational lensing near fold catastrophes. Using a generic form for the lensing map near a fold, we determine the observable properties of the lensed images, focusing on the case when the individual images are unresolved, i.e., microlensing. Allowing for images not associated with the fold, we derive analytic expressions for the photometric and astrometric behavior near a generic fold caustic. We show how this form reduces to the more familiar linear caustic, which lenses a nearby source into two images which have equal magnification, opposite parity, and are equidistant from the critical curve. In this case, the simplicity and high degree of symmetry allows for the derivation of semi-analytic expressions for the photometric and astrometric deviations in the presence of finite sources with arbitrary surface brightness profiles. We use our results to derive some basic properties of astrometric microlensing near folds, in particular we predict for finite sources with uniform and limb darkening profiles, the detailed shape of the astrometric curve as the source crosses a fold. We find that the astrometric effects of limb darkening will be difficult to detect with the currently planned accuracy of the {\\it Space Interferometry Mission (SIM)} for Galactic bulge sources; however, this also implies that astrometric measurements of other parameters, such as the size of the source, should not be compromised by an unknown amount of limb darkening. We verify our results by numerically calculating the expected astrometric shift for the photometrically well-covered Galactic binary lensing event OGLE-1999-BUL-23, finding excellent agreement with our analytic expressions. Our results can be applied to any lensing system with fold caustics, including Galactic binary lenses and quasar microlensing. ", + "introduction": "Gravitational lensing has proven to be an exceptional tool for studying a diverse set of astrophysical phenomena. Its utility is due, at least in part, to the fact it operates in a number of qualitatively different regimes. The term strong lensing, or macrolensing, is usually applied when a distant source (typically cosmological) is lensed into multiple, resolved images by an intervening mass, such as a foreground cluster or a galaxy. Weak lensing is used to refer to the case when multiple images are not created, and the gravitational field of the intervening matter serves only to slightly distort the image of the source. For most applications of both strong and weak lensing, the source, lens, and observer can be considered static. The term microlensing is often used to describe the case when multiple images are created, but are not resolved. Typically the separation of the images created by a gravitational microlens are of order the Einstein ring radius, \\begin{equation} \\thetae=\\sqrt{{{4 G M}\\over \\drel c^2} }. \\label{eqn:thetae} \\end{equation} Here $M$ is the mass of the lens, $\\drel$ is defined by, $\\drel\\equiv \\dos\\dol/\\dls$, and $\\dos$, $\\dol$, and $\\dls$ are the distances between the observer-source, observer-lens, and lens-source, respectively. In cosmological contexts, angular diameter distances should be used. When $\\thetae$ is less than the resolution, individual images in general cannot be distinguished. Due to the small scale of $\\thetae$, it is typically not a good approximation to assume that the source, lens, and observer are static. Therefore the lensing properties can be expected to change on timescales of order the Einstein ring crossing time, \\begin{equation} \\te = {\\thetae \\dol\\over v_\\perp}, \\label{eqn:te} \\end{equation} where $v_\\perp$ is the transverse speed of the lens relative to the observer-source line-of-sight. The standard observables in gravitational microlensing are therefore the time rate of change of the total magnification and center-of-light (centroid) of all the microimages. There are two different regimes where microlensing has been discussed: quasar microlensing \\citep{wambs2001} and microlensing in the Local Group \\citep{pac1996}. In the Local Group, gravitational microlensing occurs whenever a massive, compact object passes close to our line of sight to a more distant star. Microlensing was originally suggested as a method to detect baryonic dark matter in the halo of our galaxy \\citep{pac1986}, but has been developed and applied as an important tool in studying a number of astrophysical topics, including the stellar mass function \\citep{gould1996}, extrasolar planets \\citep{mandp1991}, stellar atmospheres \\citep{gould2001}, and stellar multiplicity \\citep{alcock2000,udalski2000}. The only microlensing effect currently observable is the magnification of the background source as function of time. This is because, for typical distances in the Local Group, the angular Einstein ring radius is $\\thetae\\simeq 1~{\\rm mas}({M / \\msun})^{1/2},$ and therefore too small to be resolved with current instruments. The timescale for a microlensing event is $\\te \\sim 40~\\days$. In general, it is much easier to determine the center-of-light of an image than it is to resolve it. Thus future interferometers, such as the {\\it Space Interferometry Mission (SIM)}, although still not able to resolve separations of ${\\cal O}(\\rm mas)$, should be able to measure the centroid of all the images to much better than this, perhaps even down to $10~\\muas$ in the case of {\\it SIM}. Such accuracy is sufficient to easily detect the motion of the centroid of the images created in a microlensing event, which is also of order $\\thetae$. This regime is typically referred to as astrometric microlensing, as opposed to photometric microlensing when only the total magnification is observable. Astrometric microlensing has a number of important applications. By combining ground-based photometry of microlensing events with photometry and astrometry from an astrometric satellite on an Earth-like orbit, the masses of microlenses can routinely be measured \\citep{pac1998, bsvb1998, gs1999}, allowing for the determination of the compact object mass function in the bulge, including stellar remnants \\citep{gould2000}. Astrometric information alone allows for the precise (few \\%) measurement of the masses of nearby stars \\citep{pac1995}. Finally, for a subset of events, it will be possible to obtain precision measurements of angular diameters of stars in the Galactic bulge using astrometric information \\citep{pac1998,ggh2002}. Binary microlenses have proven to be enormously useful in photometric microlensing studies. This is primarily because binary lenses exhibit caustics: closed curves on which the mapping from the lens plane to light source plane becomes critical, and the point-source magnification becomes formally infinite. Regions near caustics exhibit large, rapidly changing (with respect to source position) magnification, and are therefore useful not only for providing a large source flux, but also high angular resolution. However, binary lenses have also proven to be difficult to study both theoretically and observationally. This is partially because the lens equation, which describes the mapping from the lens plane to the light source plane, is equivalent to a fifth-order complex polynomial in the source position \\citep{witt1990}, and therefore is not analytically solvable in general. Furthermore, care must be taken when considering finite source effects near caustics due to the divergent magnification. However, considerable progress can be made when one realizes that the smooth arcs (away from cusps) of caustics that arise in nearly equal-mass binary lenses are well approximated as simple linear fold catastrophes, which have generic, simple, and most importantly, {\\it analytic} behavior. Thus the caustics of binary lenses can be analyzed analytically or semi-analytically without reference to the global (and non-analytic) topology of the general binary lens. In particular, a simple equation for the magnification of a source near a fold exists \\citep{sef92}, which has been used in a number of important applications including binary-lens fitting \\citep{albrow1999}, stellar atmospheres \\citep{gandg1999}, and caustic-crossing predictions \\citep{jandm2001}. In contrast to the astrometric behavior of single lenses, which is analytic and has been quite well studied \\citep{walker1995,jhp1999}, there have been only a few preliminary studies of the astrometric properties of binary gravitational lenses \\citep{hcc1999, ch1999,gh2000}. It is known that astrometric binary-lens curves exhibit complex behavior, including instantaneous ${\\cal O}(\\thetae)$ jumps in the image centroid trajectory that occur when a point source crosses a binary-lens caustic and two highly-magnified images appear in a position unrelated to the position of the centroid of the other three binary-lens images. The generic behavior of these centroid jumps, or how they are altered by finite source effects, is not understood. As is the case for photometric microlensing, the astrometric behavior of binary lenses will likely prove quite useful for several applications. The usefulness of binary lenses is primarily related to the complex image centroid trajectories and the large centroid jumps. Although, in general, these properties do not allow one to measure any additional parameters over the single lens case; they do allow one to measure these parameters much more easily. In particular, \\citet{gg2002} have shown that lens mass measurements can be made to a given accuracy with 1-3 orders of magnitude fewer photons with caustic-crossing binary-lens events than with single lens events, thus greatly reducing the resources required to achieve one of the primary proposed science goals of astrometric microlensing. Caustic crossing binary-lens events are also enormously useful for measuring the angular radii of microlensing source stars in the bulge, for two reasons. First, the expected ratio of binary-to-single lens events for which the source star is resolved is a factor of $\\ga 4$ for giant sources and $\\ga 10$ for main sequence sources. Furthermore, the large and complex centroid shifts expected for caustic-crossing binary-lens events makes the requisite astrometric measurements easier. \\citet{ggh2002} have shown that, by combining accurate ground-based photometry with a handful of precise astrometric measurements, caustic-crossing binary-lens events should yield $\\sim 5\\%$ stellar radius measurements with reasonable expenditure of resources.\\footnote{Although it is possible to measure stellar angular sizes using astrometric information alone, this generally requires very densely-sampled measurements, since the source is resolved only for a short time.} Thus, given the importance of caustic crossing binary-lens events, an analytic study of the generic behavior of astrometric microlensing near folds would prove quite useful. In quasar microlensing, the separate macroimages of a quasar that is multiply-imaged by a intervening galaxy or cluster also feel the combined, non-linear effect of individual point masses (i.e. stars) in the macrolensing object that are near the macroimage position. The individual macroimages are in fact composed of many, unresolved microimages with separations of order the Einstein radius of a $M\\sim \\msun$ object at cosmological distances, $\\thetae \\simeq 1 \\muas(M/\\msun)^{1/2}$. The typical timescale for the source to cross an angle of $\\thetae$ is $\\te \\simeq 15~{\\rm years}$; however microlensing light curves should show structure on much smaller timescales due to the combined effects of many individual microlenses. Since it was first discussed by \\citet{cr1979}, cosmological microlensing has been studied theoretically by numerous authors (see \\citealt{wambs2001} and references therein), and detected in at least two systems (Q2237+0305, \\citealt{irwin1989,corrigan1991,wozniak2000}, B1600+434, \\citealt{koop2000}). Observations have been used to place constraints on, e.g.,\\ the size of the emitting region of quasars \\citep{wambs1990, wyithe2000b} and the mass function of microlenses \\citep{schmidt1998,wyithe2000a, koop2000}. Quasar microlensing differs markedly from microlensing in the Local Group in that the surface mass density in units of the critical density for lensing (the ``optical depth'') is of order unity, rather than ${\\cal O}(10^{-6})$ for the Local Group. In the high optical depth regime the lensing effects of the individual microlenses add nonlinearly, resulting in a complex caustic network. Due to this nonlinear behavior and the large number of lenses typically involved, calculation of the observable properties of such a lensing system is difficult and time consuming. Although in the high optical depth regime the caustics often exhibit considerably more complicated global behavior than the caustics of binary lenses in the Local Group, it is still the case that the smooth arcs (away from cusps) of the caustics are locally well-approximated by generic fold catastrophes. This fact, combined with a simple formula for the magnification near folds, has been exploited by numerous authors to quickly and efficiently calculate various observable properties of quasar microlensing \\citep{wambs1991,lw1998,ww1999,fw1999}. The observable effects of quasar microlensing have been limited to the relative magnifications of the various macroimages as a function of time. As with Local Group microlensing, astrometric effects should also be present. The centroid of the individual macroimages should vary as a function of time, particularly when new images are created or destroyed when the source crosses a caustic. This effect has been studied by \\citet{ws1995} and \\citet{lw1998}. In particular, \\citet{lw1998} predict that magnitude of the centroid shift for the Q2237+0305 system (apparent magnitude $R\\la 18.5$) can be as large as $\\sim 50\\muas$, and thus potentially observable with {\\it SIM}. They also note that the magnitude of the centroid shift is often correlated with the magnitude of the change in total magnification. The analytic results presented in this paper may prove useful for this application. Here we study the generic, local behavior of microlensing near fold catastrophes. In \\S 2 we present an analytic study of the photometric and astrometric behavior near folds. We begin with the equations that describe the mapping near a fold caustic in \\S 2.1, and use these to derive the behavior of a point source near a fold. We extend this analysis to finite sources in \\S 2.2, and limb-darkened sources in \\S 2.3. In \\S 2.4 we show how and when our generic parabolic fold form reduces to the more familiar linear caustic. In \\S 2.5 we use our analytic results to derive some generic results about the astrometric behavior near folds. We verify the applicability of our results in \\S 3 by numerically calculating the photometric and astrometric behavior of one well-observed binary-lens event. We find excellent agreement with our analytic formulae. Finally, we summarize and conclude in \\S 4. Our goal is to provide a thorough, comprehensive study of gravitational microlensing near fold caustics. Although our study is interesting in its own right, the primary utility of the results presented here is their potential application to the topics mentioned in the previous paragraphs. A prescription for how specifically our results can be applied to these topics is beyond the scope of this paper, but we will make general comments along these lines over the course of the paper, and more specific comments in \\S 2.5. We are currently preparing a complementary, similarly detailed study of microlensing near cusps. Combined with this study, we will have a reasonably thorough and complete understanding of the local behavior of microlensing observables near all stable gravitational lensing singularities. We note that some of the results derived here, particularly the results on the photometric behavior near folds, have been presented elsewhere (see, e.g. \\citealt{sef92}, \\citealt{plw01}, and \\citealt{fw1999}). We include those results here for the sake of completeness. ", + "conclusions": "We have presented a detailed study of gravitational lensing near fold catastrophes, concentrating on the regime where the individual images are unresolved, i.e.\\ microlensing. By Taylor expanding the scalar potential $\\psi$ in the neighborhood of a fold up to third order in the image position, one can obtain a generic form for the lensing map near a fold. Beginning with this mapping, we derive the local lensing properties of a source in the vicinity of the fold caustic. Approximating the critical curve by its tangent line at the origin, we find that the caustic is locally a parabola. On one side of the parabola, the fold lenses a nearby source into two images; on the other side of the parabola, there are no images. We derive the image positions and magnifications as a function of the position of the source. We find that the magnifications of the two images are equal, and recover the well-known result that the magnification is inversely proportional to the square root of the distance to the caustic. We show how this holds for parabolic caustics (as well as linear caustics), provided that the `vertical' distance from the caustic is used. Assuming a rectilinear source trajectory, and allowing for the existence of slowly- and smoothly-varying images not associated with the fold caustic, we derive analytic expressions for the total magnification and image centroid (center-of-light) as a function of time. We then consider how the photometric and astrometric behavior is altered in the presence of a finite source size. We derive semi-analytic expressions for the magnification and centroid as a function of time for both a uniform source, and limb-darkened source. Along the way we derived expressions that can be used to evaluate the photometric and astrometric behavior near a fold for a source with arbitrary surface brightness profile. We then show how and under what conditions the generic parabolic fold reduces to the more familiar linear fold. We derive simplified expressions for the individual and total image positions and magnifications near a linear fold. We used some of our analytic results to derive a few generic properties of microlensing near folds. In particular, we derive and evaluate expressions for the magnitude of the centroid jump that occurs when a finite source crosses a fold relative to the point source jump, and the magnitude of the effect of limb darkening on both the photometric and astrometric behavior. Notably, we predict, for Galactic bulge lensing events, the shape of the centroid due to finite sources with uniform and limb darkening surface brightness profiles. We also find for Galactic bulge lensing that the effect of limb darkening on the image centroid near a fold is quite similar to the uniform source case, making the limb darkening effect difficult to detect by the currently planned accuracy for the instrumentation of {\\it SIM}. We discussed how our formulae can be used to fit both photometric and astrometric data sets near fold caustic crossings and thus used to derive such properties as the angular size of the source and the microlensing parallax. Finally, we numerically calculate expected astrometric behavior of the photometrically well-observed Galactic bulge binary lensing event {\\event} \\citep{albrow2001}, finding excellent agreement with our analytic predictions. Caustics are ubiquitous in gravitational lenses, and the most common type of caustic is the fold. Caustics play an especially important role in microlensing, as the rapid time variability of the total image magnification allows the possibility of detailed studies of the source and lens. In the future, we can expect that time-series photometric measurements will be supplemented by time-series {\\it astrometric} measurements of the center-of-light of microlens systems. This paper presents the most thorough and comprehensive study of the photometric and astrometric behavior of gravitational microlensing near fold caustics to date. The results should prove useful to those studying microlens systems with caustics: The analytic expressions derived here can be used to fit fold caustic crossings observed both photometrically and astrometrically, gain some insight into more complicated numerical studies, and establish predictions for the feasibility of future observations." + }, + "0112/astro-ph0112194_arXiv.txt": { + "abstract": "Two selection statistics are used to extract new candidate periodic variables from the epoch photometry of the Hipparcos catalogue. The primary selection criterion is a signal to noise ratio. The dependence of this statistic on the number of observations is calibrated using about 30 000 randomly permuted Hipparcos datasets. A significance level of 0.1\\% is used to extract a first batch of candidate variables. The second criterion requires that the optimal frequency be unaffected if the data are de-trended by low order polynomials. We find 2675 new candidate periodic variables, of which the majority (2082) are from the Hipparcos ``unsolved\" variables. Potential problems with the interpretation of the data (e.g. aliasing) are discussed. ", + "introduction": " ", + "conclusions": "" + }, + "0112/astro-ph0112477_arXiv.txt": { + "abstract": "We report radial velocities for 844 FGKM--type main sequence and subgiant stars and 45 K giants, most of which had either low-precision velocity measurements or none at all. These velocities differ from the standard stars of \\cite{Udry99a} by 0.035 \\kms (RMS) for the 26 FGK standard stars in common. The zero--point of our velocities differs from that of Udry et al.: $$ = +0.053 \\kmse. Thus these new velocities agree with the best known standard stars both in precision and zero--point, to well within 0.1 \\kmse. Nonetheless, both these velocities and the standards suffer from three sources of systematic error, namely, convective blueshift, gravitational redshift, and spectral type mismatch of the reference spectrum. These systematic errors are here forced to be zero for G2V stars by using the Sun as reference, with Vesta and day sky as proxies. But for spectral types departing from solar, the systematic errors reach 0.3 \\kms in the F and K stars and 0.4 \\kms in M dwarfs. Multiple spectra were obtained for all 889 stars during four years, and 782 of them exhibit velocity scatter less than 0.1 \\kmse. These stars may serve as radial velocity standards if they remain constant in velocity. We found 11 new spectroscopic binaries and report orbital parameters for them. ", + "introduction": "\\label{intro} The radial velocity of a star is ideally the component of the velocity vector of its center of mass that lies along the line--of--sight. Radial velocities are valuable for a variety of astrophysical investigations, including studies of the structure of the Milky Way Galaxy, the orbits of long--period binary stars, and the distances to star clusters (see for example, \\citet{bm98}). ``Barycentric'' radial velocities (sometimes referred to as ``absolute'' radial velocities), such as reported here, are measured relative to the barycenter, or center of mass, of the Solar System. Such velocities are often (incorrectly) termed ``heliocentric'', though the Sun moves with a speed of $\\sim$13 \\ms relative to the barycenter. Radial velocities of stars in the Galaxy are often measured with an accuracy of only $\\sim$0.5 \\kmse. With advances in the accuracy of proper motion measurements to $\\sim$1 mas/yr (e.g., Perryman et al. 1996) for many stars, a corresponding increase in the accuracy of radial velocities is required. Meanwhile, the best {\\em relative} radial velocities have precisions of 3 \\ms \\citep{But00} and new instruments (e.g. HARPS) are designed to achieve a precision of 1 \\ms \\citep{Bou01, pepe00}. These relative velocities have proved useful in the detection of extrasolar planets (e.g., \\citet{MCM00}) but they are not necessarily tied to a velocity zero--point. Nonetheless, the precision--velocity instruments have overcome many observational and technical hurdles related to spectroscopic Doppler--shift measurements, either by using a gas absorption cell or a fiber--fed comparison lamp spectrum \\citep{Val95, But96, Bar00}. The precision--velocity technology has been applied to the establishment of barycentric radial velocities, most notably by the Geneva team (Udry et al. 1999a,b). They have measured velocities for 38 stable dwarf stars with precision better than 0.05 \\kmse. Here we provide barycentric radial velocities with typical accuracies of 0.3 \\kms (and precise to 0.03 \\kms for a given spectral type) for 889 stars. Our intent is to provide a velocity measurement at the current epoch for a variety of purposes. Velocity variations with a time scale of hundreds of years may be detected by comparison of present and future velocities. We also hope to establish radial velocity standard stars, by indentifying a subset that exhibit no significant velocity variation above 10 \\mse. ", + "conclusions": "We have provided barycentric radial velocities with an internal precision of 0.03 \\kms for 889 stars. The error estimates stem both from the internal errors found from the spectral chunks within each spectrum and from the comparison with accurate velocities on the CORAVEL scale \\citep{Udry99a}. The radial velocities of the F, G and K dwarfs reside on a velocity zero--point defined by the observations of the Sun, using the day sky and Vesta as proxies. Our velocity scale differs by only 0.053 \\kms from that of \\citet{Udry99a} and 0.015 \\kms from that of \\citet{Stef99}, thus adding confidence to the zero points of all three sets of velocities. The radial velocities of the M dwarfs reside on the velocity system defined by \\citet{Mar87} and have not been further corrected, nor is such a correction known to be necessary. These M dwarf velocities are probably accurate to within 200 \\mse. The Doppler shifts reported here have such high accuracy that gravitational redshift and convective blueshift impose comparable (or greater) wavelength shifts. These effects were somewhat removed from our velocity measurements by using the Sun for the velocity zero--point. We expect therefore that for G2V stars the present velocities represent their ``true'' kinematic velocities within 0.03 \\kmse. However, for stars departing from solar--type the sum of the two astrophysical effects will produce systematic errors dependent on spectral type. From F to K type dwarfs this variation will be as large as $\\sim$0.3 \\kms and will cause our velocities to be low for F type stars and high for K type stars. Using estimates for these effects we give rough velocity corrections in equations (\\ref{early}) and (\\ref{late}). We presume that these corrections bring the velocities within $\\sim$0.15 \\kms of their ``true'' kinematic values. These precise radial velocities can be used to complement future highly precise proper motion measurement, such as those projected to be obtained by the GAIA mission of the ESA. The radial velocities and proper motions will give the three components of space motion of stars. These precise space motions may be useful for discerning the membership of moving groups, since young moving groups have velocity dispersions of $\\sim$0.5 \\kms \\citep{Jon70}. The 782 stars listed in Table 1 exhibited a velocity scatter of less than 100 \\ms during 4 years. These stars apparently exhibit relatively stable velocities on time scales of a decade and represent candidates for radial velocity standard stars. However, their integrity as velocity standard stars requires future observations to verify their stability. We expect that some of these 782 ``stable'' stars may reveal slow drifts in radial velocity on time scales longer than 10 yr. The accuracy of the present velocities offers an opportunity to detect such slow drifts by future measurements made with comparable accuracy. Such velocity variations may prove useful in identifying unseen companions at large orbital distances, i.e., over 10 AU. We found that 107 stars exhibited velocity variations of over 100 \\ms (RMS). For these stars, we have provided the RMS velocity, and also either an orbital solution or a description of the linear trends. We intend these measurements to provide dynamical constraints on the nature of the companions." + }, + "0112/astro-ph0112127_arXiv.txt": { + "abstract": "We report spectral and variability analysis of two quiescent low mass x-ray binaries (previously identified with ROSAT HRI as X5 and X7) in the globular cluster 47 Tuc, from a {\\it Chandra} ACIS-I observation. X5 demonstrates sharp eclipses with an 8.666$\\pm0.008$ hr period, as well as dips showing an increased $N_H$ column. Their thermal spectra are well-modelled by unmagnetized hydrogen atmospheres of hot neutron stars, most likely heated by transient accretion, with little to no hard power law component. ", + "introduction": "Many neutron stars (NSs) that have been observed in outburst as soft x-ray transients have also been detected in quiescence ($L_X\\sim10^{32-34}$ ergs s$^{-1}$), for example Cen X-4 and Aquila X-1 (see Campana et al. (1998a) for a review of soft x-ray transients, also known as quiescent low-mass x-ray binaries (qLMXBs). Spectral fits with a soft (kT=0.2-0.3 keV) blackbody (BB) spectrum, often requiring a hard power-law tail of photon index $\\sim2$, have been acceptable but imply an emission area of $\\sim1$ km radius, smaller than a NS surface. However, Rajagopal \\& Romani (1996) and Zavlin et al. (1996) showed that the atmosphere of a NS, due to the strong frequency dependence of free-free absorption, shifts the peak of the emitted radiation to higher frequencies; thus a blackbody fit will derive a temperature that is too high and a radius that is too small. Brown, Bildsten, and Rutledge (1998; BBR) showed that the interior of a transiently accreting neutron star is heated during accretion by pycnonuclear reactions, bringing the interior to a steady state temperature $\\sim10^8 \\langle$ \\mdot$/10^{-10} M_{\\sun}$ yr$^{-1}\\rangle^{0.4}$ K. This heating leads to an isotropic thermal luminosity between accretion episodes of roughly $L_q=6\\times10^{33}$ ergs s$^{-1} (\\langle$ \\mdot $\\rangle/10^{-10}$ \\Msun yr$^{-1}$ (BBR). Fits of simulated low magnetic-field ($B<10^{10}$ G) hydrogen atmospheres to qLMXBs have given implied $R_{\\infty}$ of roughly 13 km ($R_{\\infty}$ is the effective radius seen by a distant observer, $R_{\\infty}=R/\\sqrt{1-2GM/Rc^2}$), with a power law component in Cen X-4 and Aql X-1 suggested to be linked to continued accretion (Rutledge et al. 2001a, 2001b, Campana et al. 1998b). The deep ROSAT HRI observation of 47 Tucanae of Verbunt \\& Hasinger (1998) resolved 5 sources, including X5 and X7, in the central core, but without spectral resolution. With {\\it Chandra}'s ACIS-I instrument, we have identified 108 sources in the central 2\\arcmin$\\times$2.5\\arcmin region, and are able to conduct spectral analysis on a few dozen of these sources (Grindlay et al. 2001a). X5 has been identified with a $V=21.6$ counterpart, which shows variability and a blue color indicative of a faint accretion disk, while a tentative upper limit for the counterpart of X7 is $V\\sim23$ (Edmonds et al. 2001). Using $F_{V}=10^{-0.4V-5.43}$ (using the 5000-6000 \\AA range) gives $F_X/F_V$ ratios of 43 and $>$166 respectively, which are typical of LMXB systems rather than CVs (typically $\\sim1$). The likely qLMXBs X5 and X7 have blackbody-like x-ray spectra (see below) which show they are indeed hot NSs. ", + "conclusions": "" + }, + "0112/astro-ph0112311_arXiv.txt": { + "abstract": "{ We present \\sax LECS, MECS, and PDS spectra of eleven X-ray selected BL Lacertae objects.\\linebreak Combining these sources with the ones presented elsewhere we have a sample of 21 BL Lacs from the Einstein Medium Sensitivity and Einstein Slew Survey. The sample shows strong correlations of several physical parameters with the peak frequency of the synchrotron branch of the spectral energy distribution. In particular the peak frequency is correlated to the X-ray spectral shape: objects with the peak near to the X-ray band show harder and straighter X-ray spectra than those of the low energy peaked sources. This work shows that the recently proposed unification scenario for different types of blazars can hold also within the class of high frequency peaked BL Lac objects. ", + "introduction": "BL Lac objects mark one extreme population of active galactic nuclei (AGN). They exhibit intense variability (e.g. Wagner \\& Witzel \\cite{variability}) and polarization (e.g. K\\\"uhr \\& Schmidt \\cite{polarization}) but do not show strong emission or absorption lines in their spectrum (e.g. Stocke et al. \\cite{stocke}). In order to understand their physical nature through the emission processes involved and their relative contribution, one has to study their spectral energy distribution (SED). Determining the SED is a telescope time consuming issue. There exist some well studied bright BL Lac objects - like Mkn 421, 1ES 2344+512 and Mkn 501 - for which detailed measurements at several wavelengths are available, but to construct a global scenario for BL Lac objects it is necessary to also derive the properties of complete samples. This means investigating fainter objects too. Because of its sensitivity and wide energy band, which allows us to study simultaneously the BL Lac spectrum from 0.1 to 100 keV, the \\sax satellite is a powerful tool for this aim. Here we present a well defined sample of 21 BL Lac objects which have been studied using \\sax. In the following Section 2 we define the sample; observations and data analysis are explained in Section 3. The X-ray data are compared to earlier results by the ROSAT satellite in Section 4, and in Section 5 we examine the properties of the BL Lac SEDs. The results are discussed in the context of unified schemes for BL Lac objects in Section 6. ", + "conclusions": "The (physical) parameters of the objects studied here seem to partly correlate with the value of the peak frequency of the synchrotron branch of the SED. The X-ray selected sample considered in this work comprises objects whose synchrotron peaks are located near/in the X-ray band (high frequency peaked BL Lac, HBL; Padovani \\& Giommi \\cite{padovani}). Those of radio selected BL Lac objects tend instead to be located in the optical or infrared bands (low frequency peaked BL Lac, LBL), as is the case of sources belonging to the 1-Jy BL Lac sample (Stickel et al. \\cite{stickel}) whose peak frequencies can be as low as $\\sim 10^{13} \\rm Hz$ (Padovani et al. \\cite{padovani01}). We find a correlation between the peak frequency and the spectral slope as well as between the spectral slope and the flattening of the soft X-ray spectrum (parameterized as ``intrinsic'' absorption). These trends are consistent with the properties of the SED of HBLs: when the peak of the SED moves into the X-ray region the X-ray spectra are expected to become harder and straighter as the peak frequency rises. This behavior also accounts for the relative high X-ray luminosities observed in X-ray dominated BL Lacs (although LBL have higher total bolometric luminosities than HBL their X-ray emission corresponds to the energy band between the synchrotron and IC components, e.g. Fossati et al. \\cite{fossati}). Although within our sample it has not been possible to determine that HBL have lower bolometric luminosity (as a large fraction of the bolometric luminosity is expected to be emitted in the IC branch, which extends to energies outside the frequency range we covered), we can assume - as already mentioned - that the radio luminosity is a good tracer of the bolometric one (Fossati et al. \\cite{fossati}, Ghisellini \\cite{ghisellini99}). The lack of a visible trend between such luminosities and the peak frequencies has then to be ascribed to the small luminosity range spanned by the HBL sample in comparison with the intrinsic dispersion of the peak-luminosity relation (as expected e.g. for a distribution of the observing viewing angle which tends to produce an opposite peak-luminosity trend). If we consider this, the absence of correlation between the total luminosity in the synchrotron branch and the peak frequency will then result in higher bolometric luminosities for the LBL in comparison to the HBL. This is consistent with the fact that LBL show strong IC components compared to the synchrotron branch, while HBL emit similar luminosities in the IC and the synchrotron branch. As the sample discussed here only includes HBL objects, we would expect the X--rays to be dominated by the higher energy part of the synchrotron and/or the lower energy part of the IC component (Padovani et al. \\cite{padovani01}, W98). Interestingly, no flattening associated with the latter component is observed up to several tens of keV (as the PDS data show that the X-ray power-law can be extended up to $\\sim 100 \\rm keV$), implying that the IC seems to start emerging at frequencies $\\nu > 10^{19} \\rm Hz$. This is fully consistent with the assumed SED for HBL as inferred from the sequence scenario (Fossati et al. (\\cite{fossati}). In summary the \\sax spectral survey shows that the X-ray properties of X-ray selected BL Lac objects are in good global agreement with the unified model for blazar, which ascribes the differences among blazars mostly to the location of the peak frequencies of the synchrotron and inverse Compton spectral components (e.g. Padovani \\& Giommi \\cite{padovani}, Ghisellini et al. \\cite{ghisellini}). Even though the correlation of peak frequency with total luminosity as described by Fossati et al. (\\cite{fossati}) cannot be found in this sample, we find that the spectral properties are fully consistent with objects having a spread in the position of the peak of the synchrotron branch around the X-ray band. The sample as a whole represents only a part of the blazar population. Comparing the luminosities in the different energy bands studied here, the objects match the average SEDs of HBL as shown by Fossati et al. (\\cite{fossati}). The energy coverage provided by \\sax toward the high energies did not allow us to directly examine the link between the position of the peaks and the bolometric emitted power, as postulated by such model. Such study will be however possible with observations with sensitive instruments and good spectroscopy resolution above $\\sim 300$ keV as provided by e.g. the INTEGRAL mission, which will cover the energy range of the spectral change from the synchrotron to the IC dominated emission and possibly of the peak of the IC component." + }, + "0112/astro-ph0112061_arXiv.txt": { + "abstract": "{We report the result of an on-going survey for 22 GHz \\hho\\ maser emission towards infrared luminous galaxies. The observed galaxies were selected primarily from the IRAS bright galaxy sample. The survey has resulted in the detection of one new maser. The new maser was discovered towards the [U]LIRG/merger galaxy NGC\\,6240, which contains a LINER nucleus. This is the first detection of an \\ho maser towards this class of galaxy, they are traditionally associated with OH megamaser sources. The detected maser emission is highly redshifted ($\\sim$ 260--300 \\kmss) with respect to the adopted systemic velocity of the galaxy, and we identified no other significant emission at velocities $\\la$ $\\pm$ 500 \\km relative to the systemic velocity. The presence of high-velocity maser emission implies the possible existence of a rotating maser disk formed in the merging process. The large maser luminosity ($\\sim$ 40 \\lsun) suggests that an active galactic nucleus could be the energy source that gives rise to the water emission. Alternatively, the maser emission could be associated with the previously observed double radio source in the centre of the galaxy. Interferometric observations with high angular resolution will be able to clarify the origin of the new maser. ", + "introduction": "The discovery of the highly Doppler-shifted maser emission at 22 GHz in the 6$_{16}$--5$_{23}$ transition of \\ho in NGC\\,4258 (\\cite{naka93}), has resulted in a resurgence of single-dish surveys for extragalactic water masers. Recent maser surveys towards active galactic nuclei (AGN) have been motivated by the VLBI observations of a sub-parsec-scale \\ho maser disk rotating around a central massive object in NGC\\,4258 (\\cite{miyo95}). The detection and imaging of such \\ho $megamaser$ emission enables us to probe the kinematics and dynamical structures of parsec-scale circumnuclear disks (\\cite{mora99}). \\ho megamasers studied to date can be categorised into two types; $disk~masers$ residing in a (sub-)parsec scale disk around an active nucleus, and $non-disk~masers$ which are seen significant distances from the nucleus. The latter can also be sub-categorised into two types: jet masers associated with radio jets e.g NGC\\,1068 (\\cite{gall96}), NGC\\,1052 (\\cite{clau98}), and Mkn\\,348 (\\cite{falc00}, \\cite{peck01}, \\cite{richards01}) or a nuclear outflowing component such as that observed in the Circinus galaxy (\\cite{linc00}). Masers associated with jets appear to have different spectral characteristics than the $disk~masers$, the spectra tend to be broad (few hundred \\kmss) and featureless, whereas the $disk~masers$ tend to consist of bright, narrow components clustered in velocity groups. \\ho megamasers generally seem to trace the nuclear activity in the central parsecs of AGN. To deduce the general properties of such megamasers and those of the central parsecs of AGN, we need to increase the number of detections of water megamasers. \\indent At this time, 24 extragalactic water masers inside AGN have been discovered with single-dish surveys (e.g. \\cite{mora99}). The detection rates of new masers remain quite low, less than $\\sim$ 7\\% (e.g., \\cite{braa96, braa97, linc97}). From the presently known sample of sources we have already deduced some properties of \\ho megamasers. They prefer to lie in type 2 Seyfert AGN, suggesting that obscuring gas around an active nucleus coupled with long gain paths along an edge-on disk in the line of sight, plays a vital role in giving rise to strong maser emission. A number of \\ho megamasers are known to contain a parsec-scale radio ``core-jet'' structure which the maser can amplify in the gaseous environment. Accordingly it is reasonable to search for \\ho megamasers towards type 2 Seyfert/LINER nuclei enshrouded by circumnuclear gas around a compact continuum nucleus. \\\\ Here we report the result of an on-going single-dish survey for \\ho megamasers towards AGN in infrared (IR)-luminous galaxies, and the discovery of the new maser in the galaxy NGC\\,6240. ", + "conclusions": "We have made the first detection of \\ho megamaser emission from a [U]LIRG/merger galaxy. We have detected emission in the LINER galaxy NGC\\,6240, the maser features are observed to be by 260-300 \\km redshifted with respect to the systemic velocity and to have a luminosity of ($\\sim$ 40 \\lsun). The maser could arise from a dense circumnuclear molecular cloud on a $\\sim 100$ parsec scale or from a spatially compact maser disk inferred from the presence of the high velocity maser feature(s). The interpretation of the observed maser is not unique, though we favour the latter case. Interferometric observations of the new maser will provide an opportunity to investigate the kinematics of the ongoing merger." + }, + "0112/astro-ph0112257_arXiv.txt": { + "abstract": "With 40~ks of \\chandra\\ ACIS-S3 exposure, new information on both the starburst and QSO components of the X-ray emission of Markarian 231, an ultraluminous infrared galaxy and Broad Absorption Line QSO, has been obtained. The bulk of the X-ray luminosity is emitted from an unresolved nuclear point source, and the spectrum is remarkably hard with the majority of the flux emitted above 2~keV. Most notably, significant nuclear variability (a decrease of $\\sim45\\%$ in approximately 6 hours) at energies above 2~keV indicates that \\chandra\\ has probed within light hours of the central black hole. Though we concur with Maloney \\& Reynolds that the direct continuum is not observed, this variability coupled with the 188~eV upper limit on the equivalent width of the \\feka\\ emission line argues against the reflection-dominated model put forth by these authors based on their \\asca\\ data. Instead, we favor a model in which a small, Compton-thick absorber blocks the direct X-rays, and only indirect, scattered X-rays from multiple lines of sight can reach the observer. Extended soft, thermal emission encompasses the optical extent of the galaxy and exhibits resolved structure. An off-nuclear X-ray source with a 0.35--8.0~keV luminosity of $L_{\\rm X}=7\\times10^{39}$~erg~s$^{-1}$, consistent with the ultraluminous X-ray sources in other nearby starbursts, is detected. We also present an unpublished FOS spectrum from the \\hst\\ archive showing the broad \\CIV\\ absorption. ", + "introduction": "\\label{sec:intro} The extraordinary galaxy Markarian~231 was discovered in 1969 as part of a survey searching for galaxies with strong ultraviolet continua \\citep{Ma1969}. With the first spectrum, it was identified as unique among the Markarian sample because of deep, blue-shifted Na~D absorption evocative of broad absorption line quasi-stellar objects \\citep[BAL~QSOs;][]{ArDiEs1971}. Soon after, \\citet{AdWe1972} also identified broad absorption lines from \\CaII\\ H and K and \\HeI\\ at the same velocity as the Na~D system. In addition, they first described the long tidal tails and disturbed morphology of the galaxy. Since those early observations, \\mrk\\ has maintained its reputation as an exceptional object and continues to be a favorite target in all wavelength regimes. Years before the {\\it Infrared Astronomical Satellite} (\\iras) was launched, \\mrk\\ was known to have an infrared luminosity on a par with QSOs \\citep[$\\sim4\\times10^{12}$~\\lsun;][]{RiLo1972,RiLo1975}. \\citet{We1973} initially proposed that this energy could be non-thermal ultraviolet emission reradiated thermally by dust. Even after \\iras\\ expanded the known population of ultraluminous infrared galaxies, \\mrk\\ remains one of the most luminous and the most powerful known at $z<0.1$ \\citep[e.g.,][]{SuEtal1998}. An association of \\mrk\\ with BAL~QSOs was suggested from the start \\citep{ArDiEs1971}, but confirmation of this status required ultraviolet spectroscopy. The first published {\\it International Ultraviolet Explorer} (\\iue) spectra showed no obvious \\CIV\\ or \\MgII\\ emission or absorption, but the signal-to-noise ratio of the data was low \\citep{HuNe1987}. {\\it Hubble Space Telescope} (\\hst) ultraviolet spectra confirmed the BAL~QSO identification of \\mrk\\ with the discovery of broad \\MgII\\ and \\FeII\\ absorption \\citep{SmScAlAn1995}. Broad absorption lines are found in $\\sim10$\\% of optically selected QSOs, but \\MgII\\ BAL QSOs only comprise $\\sim10$\\% of that population. Other spectral characteristics that \\mrk\\ shares with low-ionization BAL QSOs include reddened optical and ultraviolet continua \\citep[e.g.,][]{BoEtal1977,HuNe1987}, strong optical \\FeII\\ emission \\citep[e.g.,][]{AdWe1972,BoEtal1977}, and weak \\OIII\\ emission \\citep[e.g.,][]{BoMe1992}. Though the emission of many ultraluminous infrared galaxies appears to be dominated by energetic starbursts, \\mrk\\ has been repeatedly identified as an exception \\citep[e.g.,][]{GoJoDoSa1995,KrCoThKr1997,RiEtal1999}, and many pieces of evidence point toward an accreting black hole as the major power source behind the enormous infrared luminosity (though see Downes \\& Solomon 1998\\nocite{DoSo1998} for an alternate view). The optical spectrum exhibits broad \\Ha\\ \\citep[${\\rm FWHM}=2870$~\\kms; e.g.,][]{BoMe1992} and asymmetric \\Hb\\ emission. The bright nuclear source is unresolved at radio, infrared, and optical wavelengths \\citep[e.g., Ulvestad, Wrobel, \\& Carilli 1999;][]{LaEtal1998,SuSa1999}\\nocite{UlWrCa1999}, and a parsec-scale radio jet with low apparent speeds has been resolved with Very Long Baseline Array (VLBA) observations \\citep{UlWrCa1999}. Variability on time scales of years has been observed in the radio emission \\citep{UlEtal1999} and in the Na~D and \\HeI\\ absorption lines \\citep{BoMeMoPe1991,KoDiHa1992,FoRiMc1995}. The \\iras\\ color, $F_{25}/F_{60}$ (the ratio of the flux density at 25~$\\mu$m to that at 60~$\\mu$m), is characterized as ``warm'' ($>0.25$), $F_{25}/F_{60}=0.29$ \\citep[e.g.,][]{LoHuKlCu1988}. This implies a non-thermal continuum or warm dust, more likely to result from the hard ionizing continuum of an active nucleus. In addition, the ratio of infrared luminosity to mass in H$_2$, $L_{\\rm IR}/M_{\\rm H_2}=225~L_{\\odot}~M_{\\odot}^{-1}$, is difficult to explain without a powerful QSO contributing the bulk of the infrared luminosity \\citep*{SaScSo1991}.{\\footnote{For reference, the total \\citet{SaScSo1991} sample has a median value of $L_{\\rm IR}/M_{\\rm H_2}\\approx26~L_{\\odot}~M_{\\odot}^{-1}$, and NGC~4418, the galaxy with the next highest value after \\mrk, has $L_{\\rm IR}/M_{\\rm H_2}=89~L_{\\odot}~M_{\\odot}^{-1}$.}} In \\hst\\ images, the point-like nucleus has colors inconsistent with a starburst, and $M_{B}\\sim-21.6$ with no correction for reddening \\citep{SuEtal1998}. From the shape of the optical and ultraviolet continua and infrared line ratios, the reddening is estimated to be $A_{V}\\sim2$~mag \\citep{BoEtal1977,CuRiLe1984}. Correcting for reddening then places \\mrk\\ at QSO luminosities ($M_{B}\\lesssim-24$) that can account for the infrared energy \\citep[e.g.,][]{SuEtal1998}. Though the primary power behind the incredible far-infrared luminosity of \\mrk\\ is almost certainly an active nucleus, the galaxy is also undergoing an energetic starburst. The $B-R$ colors of optical knots are indicative of active, early-type star formation \\citep{HuNe1987}, and \\hst\\ images of the nucleus clearly resolved blue, star-forming knots \\citep{SuEtal1998}. Most dramatically, in the inner kiloparsec, a nuclear ring of active star formation with a rate estimated to be $>100$~\\msun~yr$^{-1}$ has been studied with Very Large Array and VLBA observations \\citep{CaWrUl1998,TaSiUlCa1999}. Furthermore, \\citet{SaEtal1987} observed CO emission suggestive of many spiral galaxies worth of concentrated molecular gas that has since been mapped with high resolution by \\citet{BrSc1996}. The optical morphology of \\mrk\\ is irregular, and the asymmetries in the structure of the host galaxy are likely the result of a merger. Tidal tails extending more than $30\\arcsec$ to the North and South of the galaxy clearly visible in deep images further support the merger scenario. The combination of starburst and luminous QSO make \\mrk\\ a classic example of the transition from starburst to QSO in the paradigm outlined by Sanders et al. (1988)\\nocite{SaEtal1988}. X-ray studies of ultraluminous infrared galaxies are essential since they provide a direct probe not only of the environment of an active nucleus but also of the end-products of stellar evolution. \\mrk\\ is of additional interest as a BAL QSO. BAL QSOs are notoriously weak X-ray sources, and few X-ray spectra of them exist in the literature. Since the \\heao\\ and \\einstein\\ era, \\mrk\\ has been identified as anomalously X-ray weak given the implied power of its active nucleus \\citep{EaAr1988,Ri1988}. \\mrk\\ was first detected in X-rays with the \\rosat\\ Position Sensitive Proportional Counter \\citep[PSPC;][]{RiLaRo1996}, and extended emission was observed with the High Resolution Imager \\citep[HRI;][]{Tu1999}. \\asca\\ observed this source twice, and spectral analysis indicated both starburst and non-thermal emission in the X-rays \\citep{Iw1999,Tu1999}. With the most recent 100~ks \\asca\\ observation, \\citet{MaRe2000} argued that the hard X-ray spectrum is dominated by X-rays ``reflected'' from neutral or nearly neutral material, and they modeled an \\feka\\ emission line with EW$\\approx290$~eV. In this model, the obscuration to the nucleus is so severe that no direct X-rays below $\\sim10$~keV escape along the line of sight. This implies an absorption column density $\\gtrsim10^{24}$~\\cmsq\\ and suggests an intrinsic X-ray luminosity more typical of a luminous QSO, although the precise value is highly uncertain. In this paper, we present a \\chandra\\ ACIS-S3 observation of \\mrk. Previous X-ray spectral analysis of this target has been confused by the overlap of the starburst and nuclear components due to the large point spread functions (PSFs) of earlier X-ray missions. With the excellent $1\\arcsec$ spatial resolution of \\chandra, the nuclear and disk components can be spatially separated for spectral analysis. At $z=0.042$, $1\\arcsec$ is approximately 0.8 kpc for $H_0=75$~km~s$^{-1}$~Mpc$^{-1}$ and $q_0=\\onehalf$. The low Galactic column density, \\nh$=(1.03\\pm0.10)\\times10^{20}$~\\cmsq\\ \\citep*{ElLoWi1989}, also makes \\mrk\\ an excellent target for studying the soft X-rays of an active starburst galaxy. ", + "conclusions": "\\label{sec:discussion} \\subsection{Mrk~231 as a Broad Absorption Line QSO} \\subsubsection{Spectroscopic Classification} The definitive classification of Mrk~231 as a BAL~QSO has been debated in the literature for almost three decades. The canonical definition of a BAL~QSO, as published in \\citet{WeMoFoHe1991}, is based on a conservative measurement of the equivalent width of the \\CIV\\ resonance absorption line system called the balnicity index (BI). All QSOs with BI~$>0$~\\kms\\ are considered to be BAL~QSOs. Though Mrk~231 is a bright optical source, the large amount of intrinsic reddening in its spectrum renders it relatively faint in the ultraviolet, and a high-quality spectrum of the \\CIV\\ region is not available in the literature. A search of the \\hst\\ archive revealed a Faint Object Spectrograph (FOS) calibration spectrum of \\mrk\\ taken on 1996 Nov 21 covering the crucial \\CIV\\ region. This spectrum, displayed in Figure~\\ref{fig:fos_spec} and unpublished to our knowledge, clearly shows absorption blue-shifted from the \\CIV\\ and \\NV\\ emission lines, as well as possible \\AlIII\\ and \\SIV\\ broad absorption. A rough measure of the balnicity index, BI~$\\gtrsim600$~\\kms, is a quantitative indicator of the BAL~QSO nature of this object. In addition, Mrk~231 fits within the BAL~QSO class based on its X-ray properties. The quantity \\aox, the spectral index of a power law defined by the flux densities at rest-frame 3000~\\AA\\ and 2~keV, is a useful parameter for measuring the X-ray power of a QSO relative to its ultraviolet continuum emission. A large, negative \\aox\\ indicates relatively weak soft X-ray emission; the mean value of \\aox\\ for radio-quiet QSOs is $\\approx-1.48$ \\citep[e.g.,][]{LaEtal1997} with a typical range from $-1.7$ to $-1.3$ for objects that do not suffer from X-ray absorption \\citep*[e.g.,][]{BrLaWi2000}. Weakness in soft X-rays is plausibly explained by intrinsic X-ray absorption, which strongly depresses the observed flux at low-to-moderate energies. A strong correlation of large, negative values of \\aox\\ with the absorption-line equivalent width of \\CIV\\ supports this hypothesis \\citep{BrLaWi2000}. As expected, BAL~QSOs populate the extreme end of this correlation: they are the weakest soft X-ray sources as well as the QSOs with the most extreme ultraviolet absorption. Recent spectroscopic observations of three AGN with large, negative \\aox\\ values and strong \\CIV\\ absorption, PG~1411+442, PG~1535+547, and the BAL~QSO PG~2112+059, found direct evidence of intrinsic X-ray absorption \\citep{BrWaMaYu1999,GaEtal2001a}. For comparison with the BAL~QSOs of the \\citet{BrLaWi2000} sample, we have measured the equivalent width of the \\CIV\\ absorption blueward of the expected location of the emission line, EW~$\\approx10$~\\AA. Calculating \\aox\\ to complete the comparison is a difficult matter in this complicated object. Though the rest-frame 2~keV flux density can be determined from the best-fitting X-ray model, the ultraviolet spectrum suffers from severe extinction. A na{\\\"{\\i}}ve measurement of \\aox\\ from the observed 3000~\\AA\\ flux density yields \\aox$=-1.72$; however, this can only be considered an upper bound on the value. Correcting for $A_{V}=2$~mag \\citep[e.g.,][]{BoEtal1977} using the extinction curves of \\citet{SpFo1992} results in \\aox$=-2.08$; this value includes both starburst and scattered flux as well as the reddened nuclear continuum before the reddening correction and is therefore a reasonable lower limit to \\aox. Regardless of the exact value of \\aox, the strong \\CIV\\ absorption coupled with weakness in soft X-rays places Mrk~231 at the heavily absorbed end of the \\citet{BrLaWi2000} correlation, which is populated primarily by BAL~QSOs.{\\footnote{Though \\cite{BrLaWi2000} did not correct for intrinsic reddening in their ultraviolet measurements, several of the BAL~QSOs in their sample, e.g., PG~2112+059, do not have reddened continua. Those that do, such as PG~1700+518, would only have more negative values of \\aox.}} \\subsubsection{X-ray Variability and Spectral Results} \\label{sec:discussion_var} With this \\chandra\\ observation, we can demonstrate convincingly that the 2--8~keV X-ray luminosity is dominated by the unresolved active nucleus in Mrk~231. The hard-band radial profile is consistent with the \\chandra\\ PSF, and the significant, hard-band variability provides compelling evidence that \\chandra\\ is probing within light hours of the central black hole. At the observed rate the flux would decrease by a factor of 2 in $\\approx7$~hr. The $45\\pm14\\%$ decrease in count rate in the hard band represents a luminosity difference, $\\Delta L_{\\rm X}\\approx10^{42}$~\\lumin; variability of this magnitude can only reasonably originate in the immediate vicinity of the active nucleus. In addition, the compact spatial extent and very hard spectrum of the dominant energy source preclude a significant contribution by thermal starburst emission. As an alternative, a concentrated population of ultraluminous X-ray binaries cannot reasonably reproduce the variability, the spectrum, or the luminosity in the 2.0--8.0~keV band, $L_{2-8}=10^{42.2}$~\\lumin. The measured range of \\aox, $-2.08<$~\\aox~$<-1.72$, suggests a source that is under-luminous in the X-ray band by a factor of 6--36, and absorption remains the most likely explanation for this faintness. Though the moderate signal-to-noise ratio of these data is not sufficient to convincingly constrain the multi-component spectral properties of \\mrk, progress has been made. Assuming a single, direct, highly absorbed continuum as the sole source of the nuclear X-rays would require a bizarre and contrived absorber in order to force the model to match the observed flatness of the X-ray spectrum (see $\\S$\\ref{sec:basic_spec}). Though a reflection-dominated model such as that examined by \\citet{MaRe2000} remains feasible based on the spectral fitting alone, our detection of significant variability is inconsistent with their physical picture. The observed X-ray variability during the \\chandra\\ observation strongly suggests that continuum emitted on small scales is contributing significantly to the hard-band flux. This continuum cannot reasonably be considered to be direct as the X-ray flux is not sufficient for this luminous QSO, and so we prefer an indirect (i.e., reflected or scattered) origin for this variable component. For a physically reasonable reflecting medium, the size scale is too large to account for the short observed time scale of the variability. The smallest possible source of a reflecting medium would be the BAL wind \\citep[e.g.,][]{El2000} at distances typical of the Broad Line Region \\citep[e.g.,][]{OglePhD}. Even if the BAL wind resides at the smallest radii modeled \\citep[$R\\sim10^{16}$~cm; e.g.,][]{MuChGrVo1995}, however, the shortest expected time scale of variability for the reflected continuum would be days to weeks as opposed to the hours observed. Therefore, even assuming the smallest reasonable scale for reflecting material, a reflection-dominated continuum is not favored by our observation of hard-band variability. Our alternate model for the nuclear continuum is comprised of three scattered power-law spectra absorbed by different column densities of gas, and this model also provides a statistically acceptable fit to the data. We interpret these three power-law components as scattered rather than direct based on luminosity arguments; after correcting for absorption, none of them could provide the 2--10~keV X-ray luminosity of $L_{\\rm X}\\gtrsim10^{44}$~\\lumin\\ expected from the infrared power \\citep[e.g.,][]{RiGiMaSa2000}. The short time scale of the variability provides a spatial scale for at least one of the scattering regions of $R\\lesssim10^{15}$~cm. An electron mirror Thomson scattering X-rays would require densities of $n\\gtrsim10^{9}$~\\cc\\ on these small scales (assuming a Thomson optical depth of $0.5$). Given the implied X-ray luminosity, particle density, and size, the ionization state of this gas can be estimated: $\\xi=L_{\\rm X}/(nR^{2})\\sim10^{5}$~erg~cm~s$^{-1}$. This scenario supplies a coherent explanation for the hard-band variability as well as the lack of detected \\feka\\ emission. The variability above $\\sim2$~keV can be explained by flux changes of a single component, a scattered power law with intrinsic absorption of \\nh$=10^{22.5}$~\\cmsq. Strong \\feka\\ line emission would not be expected from scattering off of a highly ionized plasma ($\\xi\\sim10^{5}$~erg~cm~s$^{-1}$) where the Fe atoms are largely stripped of electrons \\citep{KaBa2001}. Our lack of detection of any direct component to the top of the \\chandra\\ band-pass supports the conclusion of \\citet{MaRe2000} that an intrinsic absorber with a Compton-thick column density of $\\gtrsim10^{24}$~\\cmsq\\ blocks these X-rays. Furthermore, for a reasonable geometry, the blockage of the direct emission combined with the presence of the scattered emission implies that the Compton-thick absorber is on a roughly similar scale to the scattering mirror; absorbers much larger than the mirror would see it as a point source and would thus block it entirely (see Figure~\\ref{fig:model} for a possible structure). Given the small scales implied for this absorber as well as the requisite column density, this absorber could be identified with the ``hitchhiking gas'' proposed by \\citet{MuChGrVo1995}, which shields the BAL wind from becoming completely ionized by extreme ultraviolet and soft X-ray photons. The hitchhiking gas would not contribute a reflected continuum because it is so highly ionized, though it could scatter X-rays into the line of sight. Thus the hitchhiking gas could provide both the Compton-thick absorber along the direct line of sight and the small-scale scattering medium that preserves short timescale variability. For the absorbers of the multiple scattered power laws, the BAL wind is a reasonable candidate. As mentioned above, radiative acceleration models put the BAL wind on much larger scales ($\\gtrsim10^{16}$~\\cmsq) than our proposed scatterer, and the column densities for this wind are expected to be $\\sim10^{23}$~\\cmsq\\ \\citep{PrStKa2000}. This column density is similar to the X-ray column densities measured for BAL~QSOs from X-ray spectroscopy \\citep{MaMaGrEl2001,GrEtal2001,GaBrChGa2001b}, as well as the column densities of two of the three scattered power-law components in our modeling. Multi-epoch observations of \\mrk\\ are certainly required to investigate the possibility of multiple scattered components further. These individual absorbed power-law components might be expected to vary with time delays dependent on the physical locations of their scattering media; in this case, flux variations in the different energy bands $<1.5$~keV, 1.5--5~keV, and $>5$~keV would provide additional support for this proposed model. Based on \\rosat\\ data, \\citet*{BoBrFi1996} and \\citet{LaEtal1997} found a significant correlation between the X-ray photon indices of radio-quiet AGN and the value for the ${\\rm FWHM}_{{\\rm H}\\beta}$, where those AGN with the narrowest broad Balmer lines, Narrow-Line Seyfert~1 galaxies (NLS1s), tended toward the steep end of the distribution of $\\Gamma$. Subsequent work with \\asca\\ in the 2--10~keV band has extended this finding to higher energy, although the spread of $\\Gamma$ in the 2--10~keV band is much smaller \\citep[e.g., Brandt, Mathur, \\& Elvis 1997;][]{ReTu2000}.\\nocite{BrMaEl1997} Mrk~231 has relatively narrow Balmer emission lines, ${\\rm FWHM}_{{\\rm H}\\alpha}=2870$~\\kms\\ \\citep{BoMe1992} and ${\\rm FWHM}_{{\\rm H}\\beta}=3000$~\\kms\\ \\citep{LiTeMa1993}, particularly for a QSO of its luminosity. As a low-ionization BAL~QSO, Mrk~231 shares several spectral characteristics with NLS1s such as weak \\OIII\\ lines and strong optical \\FeII\\ emission. This similarity suggests a physical link between the low-ionization BAL~QSOs and the NLS1s; perhaps both populations are radiating at a higher fraction of the Eddington luminosity \\citep[e.g.,][]{BrGa2000}. In order to pursue this potential link, a direct measurement of the underlying photon index of a low-ionization BAL~QSO would offer powerful evidence for a connection. Unfortunately, with these data and the evident complexity of the X-ray spectrum, we are not able to constrain the underlying photon index of the X-ray spectrum. The extreme X-ray faintness of other low-ionization BAL~QSOs \\citep[e.g.,][]{GaEtal1999,GrEtal2001} may prevent a direct measurement of $\\Gamma$ in these objects for many years. The complexity we have found in the spectrum of \\mrk\\ should be remembered when interpreting X-ray data on more distant (and therefore fainter and unresolved) low-ionization BAL~QSOs. \\subsection{Mrk~231 as a Starburst Galaxy} In this observation of Mrk~231, \\chandra\\ has resolved the extent and structure of an X-ray luminous starburst. Over the entire galaxy, the 0.5--2.0~keV luminosity, $L_{0.5-2.0}$, equals $10^{41.2}$~\\lumin, with $\\sim30\\%$ from the inner $2\\arcsec$ (1.6~kpc). The strong Fe~L shell emission of the nuclear spectrum provides evidence for a $\\sim1$~keV thermal component in addition to the non-thermal spectrum. This thermal X-ray component supports the radio evidence for star formation in the subkiloparsec gas disk \\citep{CaWrUl1998}. In this disk, the star-formation rate claimed by \\citet{TaSiUlCa1999} based on their radio observations is 220~\\msun~yr$^{-1}$. Such an active starburst region might be expected to be more X-ray luminous; M82 has an 0.5--2.0~keV luminosity of $\\sim10^{40.7}$~\\lumin\\ with a star-formation rate an order of magnitude lower \\citep[e.g.,][]{PtaEtal1997}. Though its total flux from star formation places Mrk~231 at the X-ray luminous end of the starburst population, it is still about an order of magnitude fainter from 0.5--2.0~keV than the most X-ray luminous starburst, NGC~3256 \\citep{MoLeHe1999}. The starburst emission certainly has more than one component, though discriminating between a two-temperature thermal plasma and a thermal plasma plus power-law model is not possible at present. The temperature of the thermal plasma within the nuclear region is $\\sim1$~keV, and the data are consistent with such a component contributing at all radii in the galaxy. However, the best-fitting plasma temperature at a given radius should only be considered an emission-measure-weighted average as the actual X-ray spectrum is likely to be complex. Higher counting statistics and better low-energy calibration of ACIS-S3 below 0.5~keV would contribute significantly to understanding the distribution of hot gas in this starburst in more detail. The one possible off-nuclear source is consistent with the luminous point sources seen in other starburst galaxies. One source hardly characterizes a population, however, and further observations would allow a more definitive description of the luminous X-ray binary population of this galaxy." + }, + "0112/astro-ph0112468_arXiv.txt": { + "abstract": "Observations of the transient accreting pulsar \\srcnm\\ made with the \\emph{Rossi X-ray Timing Explorer} during the course of the 1998 September--November outburst, reveal a cyclotron resonance scattering feature (or ``cyclotron line'') in the hard \\xray\\ spectrum near 35\\,keV. We determine a centroid energy of $36.2_{-0.7}^{+0.5}$\\,keV, which implies a magnetic field strength of $3.1(1+z) \\times 10^{12}$\\,G, where $z$ is the gravitational redshift of the scattering region. The optical depth, $\\tau = 0.33_{-0.06}^{+0.07}$, and width, $\\sigma = 3.37_{-0.75}^{+0.92}$\\,keV, are typical of known cyclotron lines in other pulsars. This discovery makes \\srcnm\\ one of thirteen pulsars with securely detected cyclotron lines resulting in direct magnetic field measurements. ", + "introduction": "The transient accreting \\xray\\ pulsar \\srcnm\\ was discovered during an outburst in 1998 September with the All Sky Monitor (ASM) on board the \\emph{Rossi X-ray Timing Explorer} (\\rxte) \\citep{Smi98}. At the same time, 15.8 second pulsations were observed with the \\cgro/BATSE, with the pulsating source designated GRO~J1944+26 \\citep{Wil98}. The best \\xray\\ position \\citep[\\sax/MECS: 1\\amin\\ radius, 90\\% confidence, ][]{Cam98} falls within the {\\textit{Ariel V}} error region for the transient source 3A~1946+274, making it likely that the two objects are identical \\citep{Cam99}. The initial outburst (see Fig.~\\ref{f:asmlc}), which reached \\aprx6\\,cps in the \\rxte/ASM, lasted about 100\\,days and was followed by a series of smaller flares. \\citet{Cam99} discussed the extended \\rxte/ASM lightcurve, finding that the flaring was nearly periodic with a repetition period of \\aprx80\\,d. They interpret this as either the half or full binary orbital period. Either case is consistent with the known orbital period range of Be/\\xray\\ binary pulsars. Given its probable orbital period and the outburst characteristics, \\srcnm\\ is most likely another example of a Be star/\\xray\\ binary pulsar transient. This class of binaries accounts for over half of the known accreting pulsars \\citep{Liu00}. Several \\rxte\\ observations (see Figure~\\ref{f:asmlc}) were made spanning the peak of the initial outburst. In this \\emph{Letter}, we report on the spectral analysis of these observations and, in particular, the discovery of a cyclotron resonance scattering feature, or ``cyclotron line'', at \\aprx35\\,keV. \\citet{San01} report their independent discovery of the \\srcnm\\ cyclotron line with \\sax\\ in this volume. Cyclotron lines result from the scattering of \\xrays\\ by electrons in quantized Landau orbits in the \\aprx10$^{12}$\\,G fields near the magnetic poles of accretion powered pulsars. The characteristic energy of the Landau transition scales with the magnetic field as $E_{cyc} = (11.6\\text{\\,keV})(1+z)^{-1}B_{12}$, where $E_{cyc}$ is the cyclotron line energy in keV, $z$ is the gravitational redshift in the scattering region, and $B_{12}$ is the magnetic field in units of $10^{12}$\\,G. Because of this proportionality, cyclotron lines give us the only direct measure of the neutron star magnetic field. In general, harmonically spaced lines (corresponding to higher order Landau transitions) may exist \\citep[see for example][]{Hei99,Hei00,San99}. Depending on the temperature and geometry of the emitting and scattering material and the viewing angle with respect to the magnetic field, the line profiles may be broad and complex \\citep[e.g.][]{Ara00,Kre00}. To date, about a dozen pulsars have well-established cyclotron lines, for the most part discovered with \\ginga, \\rxte, and \\sax\\ \\citep[see for example][]{Mak99,Hei00,Dal00}. ", + "conclusions": "Figure~\\ref{f:spec} shows the best fit model and residuals together with the inferred incident photon spectrum. It also shows residuals to the best NPEX model without a cyclotron line, which is then apparent in the residuals around 35\\,keV. From the cyclotron line energy, we deduce a magnetic field strength in the scattering region of $3.1(1+z) \\times 10^{12}$\\,G. The line is weakly resolved with the HEXTE energy resolution of \\aprx8\\,keV (FWHM, at 35\\,keV), and the fitted width of the line is $\\rm 9.3^{+2.5}_{-2.1}$\\% of the line energy. This is within the range of typical values determined with the Gaussian absorption profile model \\citep{Cob01b}. No harmonic line near 70\\,keV was found, but as is evident in Figure~\\ref{f:spec}, the falling high energy continuum provides inadequate statistics to make a sensitive search for such a feature. \\vspace{2ex} \\includegraphics[width=3.25in]{fig4.ps} \\figcaption{Spectral cut-off energy as a function of cyclotron line energy. A possible saturation of the correlation is seen at energies above 30\\,keV. All points are from \\rxte\\ except A0535+26 which is from HEXE/TTM \\citep{Ken94}. The line is a power law with $E_{cut} \\propto E_{cyc}^{0.7}$ \\citep{Mak99}. See text for parameter definitions.}\\label{f:ecutvsecyc} \\vspace{1ex} Figure~\\ref{f:ecutvsecyc} \\citep[after][]{Cob01b} shows the spectral cut-off energy plotted against the cyclotron line energy for 8 pulsars measured with \\rxte. Also shown is the (somewhat controversial) HEXE/TTM result for A0535+26 \\citep{Ken94}. For uniformity, the pulse-phase average spectra of all sources were fit with the PLCUT continuum: \\vspace{0.5ex} \\begin{equation} \\text{PLCUT}(E) \\propto E^{-\\Gamma} \\times \\left\\{ \\begin{array}{ll} 1 & \\text{if~} E \\leq E_{\\text{cut}}\\text{;} \\\\ e^{-(E-E_{\\text{cut}})/E_{\\text{fold}}} & \\text{otherwise} \\end{array} \\right. \\end{equation} smoothed at $E_{cut}$ \\citep[see][~for details]{Cob01b}. Two sources are excluded from the plot: 4U~1626-67, whose cut-off energy varies by a factor of four with pulse phase \\citep{Cob01b}, and the low luminosity (\\aprx$4\\times 10^{34}$\\,\\lumin) object 4U~0352+309 (X~Per) whose spectrum is not well fit by the usual pulsar models. Due to the complex shape of the 4U~0115+63 fundamental line \\citep{Hei00}, its plotted energy is half of the first harmonic value. \\begin{table*}[b] \\caption{Best fit spectral parameters for an NPEX times Gaussian absorption profile model.}\\label{t:fit} \\begin{minipage}{\\linewidth} \\renewcommand{\\thefootnote}{\\thempfootnote} \\begin{tabular}{ll} \\hline \\hline Parameter & value \\\\ \\hline Continuum & \\\\ \\cline{1-1} $\\Gamma_1$ \t\t& $2.80_{-0.27}^{+0.14}$ \\\\ $\\Gamma_2$ \t\t& $-1.38_{-0.10}^{+0.27}$ \\\\ $\\alpha$ \t\t& $(3.1_{-1.8}^{+3.2}) \\times 10^{-4}$ \\\\ $E_{fold}$ (keV) \t& $5.54_{-0.13}^{+0.16}$ \\\\ Flux\\footnote{\\eflux, 2--10\\,keV}\t& $5.5 \\times 10^{-9}$ \\\\ \\hline Cyclotron line \t\t& \\\\ \\cline{1-1} $\\rm \\tau_{cyc} $ \t\t& $0.33_{-0.06}^{+0.07}$ \\\\ $\\rm E_{cyc} $ (keV) \t& $36.2_{-0.7}^{+0.5}$ \\\\ $\\sigma_{cyc}$ (keV) \t& $3.37_{-0.75}^{+0.92}$ \\\\ \\hline $\\chi^2_r$/degrees of freedom & 0.97/49 \\\\ \\hline \\hline \\end{tabular} \\end{minipage} \\end{table*} A clear correlation exists between the line energy and spectral cut-off indicating that the cut-off is related to the magnetic field. \\citet{Mak92} and \\citet{Mak99} first noted this relationship by plotting cyclotron line energies from \\ginga\\ and other instruments, derived with a variety of continuum models, against the PLCUT cut-off energy from sometimes non-contemporaneous \\ginga\\ observations. They found that the relationship was consistent with a power law, $E_{cyc} \\propto E_{cut}^{1.4}$ (or, $E_{cut} \\propto E_{cyc}^{0.7}$), indicating a saturation as compared to a linear correlation. Figure~\\ref{f:ecutvsecyc} has the advantage that each point is derived from a uniform model fit to a single spectrum. And, with the exception of A0535+26, all points are from the same set of instruments. \\srcnm\\ fits nicely on the correlation, and below 30\\,keV, the slope is consistent with $E_{cut} \\propto E_{cyc}^{0.7}$. However, there appears to be a break in the relationship above \\aprx30\\,keV. This flattening is more abrupt than the smooth turnover of the $E_{cyc}^{0.7}$ power law suggesting that the processes that form the continuum saturate at higher magnetic fields. With \\srcnm, we have added a 13$^{th}$ accreting pulsar to the list of objects with secure cyclotron line detections. Nearly all of these have been confirmed, and several were discovered, with \\rxte\\ and \\sax. We have now begun detailed studies of these objects as a class. In particular, since ten of the thirteen sources have been observed with \\rxte, we are applying uniform analyses to all these objects to further understand the line forming regions. First results of such studies, including correlations between other line parameters, are given in \\citet{Cob01b}." + }, + "0112/astro-ph0112142_arXiv.txt": { + "abstract": "The system allowing a user to operate at server with simulated curves of spectral energy distributions (SED) and to estimate ages and redshifts by photometric data {\\bf sed.sao.ru} is described. ", + "introduction": " ", + "conclusions": "" + }, + "0112/astro-ph0112374_arXiv.txt": { + "abstract": "Intermediate resolution phase-resolved spectra of WZ Sge were obtained on five consecutive nights (July 23 -- 27) covering the initial stage of the 2001 superoutburst. Double-peaked emission lines of He\\textsc{II} at 4686 \\AA, which were absent on July 23, emerged on July 24 together with emission lines of C\\textsc{III} / N\\textsc{III} Bowen blend. Analyses of the He\\textsc{II} emission lines using the Doppler tomography revealed an asymmetric spiral structure on the accretion disk. This finding demonstrates that spiral shocks with a very short orbital period can arise during the initial stage of an outburst and may be present in all SU UMa stars. ", + "introduction": "Dwarf novae are a subclass of cataclysmic variables (CVs), which are close binary systems consisting of a white dwarf and a red dwarf secondary, transferring matter via the Roche lobe overflow (for a recent review, see \\cite{war95}). Dwarf novae show frequent outburst, which are episodes of enhanced accretion through the disk onto the central object. These systems provide good test cases for various theories of accretion disk models. WZ Sge is the prototype star of WZ Sge-type stars (originally proposed by \\cite{bai79}), which is a subclass of SU UMa-type dwarf novae. The remarkable outburst properties (a large amplitude $\\sim$ 8 mag., and the extremely long time interval between outbursts about 33 years) show that WZ Sge is the most extreme case among dwarf novae. WZ Sge has a very short orbital period of 1.37 h, and it is one of the few eclipsing dwarf novae, whose inclination angle of the orbital plane is known to be 75 $\\pm$ 2 degrees (for the detailed model of WZ Sge, see \\cite{sma93}). One of the main long-standing problems concerning accretion disks is the mechanism of angular momentum transport. Spiral shocks have long been proposed as a possible mechanism for the transport (e.g. \\cite{saw86}, \\cite{spr87}, \\cite{sav94}). If the disk extends far enough in the Roche lobe, spiral shocks are excited by the tidal field of the secondary and result in the formation of two prominent spiral arms. The gas in the disk loses its angular momentum when it passes through the spiral shocks. Analysis of the double-peaked emission lines using Doppler tomography is a well-established method for imaging the accretion disks in CVs (\\cite{mar88}). Emission lines seen in the spectra arise from the accretion disk and can be seen in a double-peaked shape as a result of Doppler motion of the disk within the binary. The velocity of the material in the disk determines the line profile. Recently, spiral structures in the accretion disks of some dwarf novae during outbursts have been discovered, IP Peg (\\cite{ste97}, \\cite{har99}, \\cite{mor00}), EX Dra (\\cite{joe00}), and U Gem (\\cite{gro01}). Their presence probed by the Doppler tomography method triggered a renewed interest in spiral shock models. The spiral pattern is interpreted as evidence for shock waves, which is consistent with the results of hydrodynamical simulations (e.g. \\cite{ste99}, \\cite{mak00}). Aside from during outbursts, \\citet{ski00} presented an extensive set of Doppler maps of WZ Sge in quiescence using both optical and infrared emission lines. In these maps, the accretion disk structure was found to be asymmetric and the bright spot region was shown to be extended along the mass transfer stream, but these structures didn't look like spirals. The outburst of WZ Sge was discovered by an amateur astronomer on 2001 July 23.565 UT at a visual magnitude 9.7 (\\cite{ish01iauc7669}), having occurred 10 years earlier than the common expectation. Receiving the report of the outburst, we started time-series spectroscopic observations, to clarify the evolution of the accretion disk structure at the very beginning of the outburst. ", + "conclusions": "The strong emmisivity of He\\textsc{II} and almost no contibution of H$\\alpha$ indicate that the accretion disk had a high temperature of a few tens of thousands K even at the edge region at the early phase of this outburst. In addition, no evidence of the secondary in the HeII and Halpha indicates that the effect of irradiation on the secondary was relatively weak. This apparent lack of strong irradiation on the secondary may have been a result of the shielding of photons by a thickened accretion disk around the superoutburst maximum, which is a natural consequence of an extremely hot (a few tens of thousands K) disk. All of the three systems in which spiral arms have been observed, i.e. IP Peg, EX Dra and U Gem, are the systems above the period gap. \\citet{har99} suggested the possibility that the spiral shocks only develop in systems above the period gap, since higher mass-ratio binaries are expected to have a heavier secondary star which induces stronger tidal torques. Nevertheless, WZ Sge is the system below the period gap, together with other SU UMa-type stars. The present discovery suggests the possibility of the existence of spirals in other SU UMa-type stars during outburst, although there has been a report of nagative detection (e.g. OY Car in outburst (\\cite{har96})). There may be a selection effect since the outburst of the short period systems are difficult to expect, and as a result, phase-resolved spectroscopy of the short period systems at the very beginning of outburst must be difficult. Further strategic observations are strongly encourged to confirm the speculation. According to \\citet{ish01iauc7669}, the humps with the orbital period called ``early superhumps'' were observed at the very early stage of this superoutburst, which is the characteristics of WZ Sge-type stars. Even though the origin of ``early superhumps'' is still unknown, the immediate evolution of spirals suggests that the double peak of ``early superhumps'' may be understood as a reflection of the two arms of the spiral structure. If so, the extended wing in the right half region of He\\textsc{II} Doppler map corresponds to the secondary maxima of the early superhumps ($\\phi \\sim 0.2$), and the weak component in left quadrant of He\\textsc{II} Doppler map corresponds to the primary maxima of the early superhumps ($\\phi \\sim 0.6-0.7$). This intensity inversion is left to be a problem. Although there is no observation of periodic modulation like ``early superhumps'' in the systems where spiral structures were detected, it may be possible that different mechanisms drive two-armed spiral patterns during outburst in WZ Sge and in some dwarf novae above the period gap. In a theoretical work, \\citet{lin79} suggested that only an accretion disk in a small mass-ratio system (such as $q<0.1$) could show a strong spiral dissipation pattern with the 2:1 Lindblad resonance. \\citet{odo90} pointed out that complex oscillations observed in AM CVn stars with more extremely low mass-ratios than WZ Sge may be understandable with multiple standing shocks demonstrated by \\citet{lin86}. We suppose that such mechanisms of standing shock of AM CVn stars may be applicable in the case of WZ Sge. \\bigskip We are grateful to Tomohito Ohshima for his detection and early notification of the long-awaited superoutburst. We are also grateful to Prof. Masanori Iye, director of OAO, who permitted us promptly to carried out the Time-Of-Opportunity observations using the OAO 91-cm telescope. We wish to thank Dr. H. C. Spruit for the use of his DOPMAP programs. We acknowledge Prof. Yoji Osaki for the useful discussion." + }, + "0112/astro-ph0112004_arXiv.txt": { + "abstract": "G347.3-0.5 (RX J1713.7-3946) is a member of the new class of shell-type Galactic supernova remnants (SNRs) that feature non-thermal components to their X-ray emission. We have analyzed the X-ray spectrum of this SNR over a broad energy range (0.5 to 30 keV) using archived data from observations made with two satellites, the R\\\"{o}ntgensatellit ($\\it{ROSAT}$) and the Advanced Satellite for Cosmology and Astrophysics ($\\it{ASCA}$), along with data from our own observations made with the Rossi X-ray Timing Explorer ($\\it{RXTE}$). Using a combination of the models EQUIL and SRCUT to fit thermal and non-thermal emission, respectively, from this SNR, we find evidence for a modest thermal component to G347.3-0.5's diffuse emission with a corresponding energy of $\\it{kT}$ $\\approx$ 1.4 keV. We also obtain an estimate of 70 TeV for the maximum energy of the cosmic-ray electrons that have been accelerated by this SNR. ", + "introduction": "The X-ray luminous Galactic supernova remnant (SNR) G347.3-0.5 (RX J1713.7-3946) was discovered during the R\\\"{o}ntgensatellit ($\\it{ROSAT}$) All-Sky Survey (Pfeffermann \\& Aschenbach 1996). Subsequent studies of this SNR's X-ray properties (Koyama et al. 1997; Slane et al. 1999) revealed that most of its X-ray emission is non-thermal and very likely produced by the synchrotron process. G347.3-0.5 therefore becomes a member of a new class of young shell-type SNRs that feature non-thermal components to their X-ray emission. Other members of this class include Cas A (Allen et al. 1997), SN 1006 (Koyama et al. 1995, Allen et al. 2001) and G266.2-1.2 (Slane et al. 2001). TeV gamma rays have been detected from the X-ray luminous northwestern rim of this SNR (Muraishi et al. 2000), making G347.3-0.5 only the third SNR (besides SN 1006 and Cas A) where such high-energy emission is detected. The presence of this emission suggests that acceleration of cosmic-ray electrons is taking place along this rim, and that additional study of this SNR's X-ray emission may lead to new insights on how SNRs act as cosmic-ray particle accelerators. \\par In order to analyze the X-ray properties of G347.3-0.5 in more detail, as well as to study how cosmic-ray particles are accelerated by this SNR, we observed this source using the Rossi X-ray Timing Explorer ($\\it{RXTE}$). We supplemented the data from these observations with publicly available data from observations that were made of G347.3-0.5 by two other X-ray satellites, $\\it{ROSAT}$ and the Advanced Satellite for Cosmology and Astrophysics ($\\it{ASCA}$). Parameters for the X-ray observations used in this analysis are listed in Table 1. By combining the data from all three satellites, we have sampled the X-ray emission from this SNR over the energy range of 0.5 through 30 keV. \\begin{table} \\caption{Summary of X-ray Observations of G347.3-0.5} \\begin{tabular}{lccccc} \\tableline & & Observed & RA & Dec & Exposure \\\\ & & Portion of & (J2000.0) & (J2000.0) & Time \\\\ Satellite & Instrument & G347.3-0.5 & (h m s) & (\\deg~\\arcmin~\\arcsec) & (Seconds)\\\\ \\tableline {\\it ROSAT} & {\\it PSPC} & All & 17 13 33.60 & $-$39 48 36.0 & 2758\\\\ {\\it ASCA} & {\\it GIS} & NW Rim & 17 12 17.76 & $-$39 35 30.8 & 20337\\\\ & & SW Rim & 17 12 53.76 & $-$39 54 23.4 & 18562\\\\ & & NE Region & 17 14 28.32 & $-$39 35 26.2 & 16220\\\\ & & SE Region & 17 15 41.52 & $-$40 02 25.8 & 40153\\\\ {\\it RXTE} & {\\it PCA} & All & 17 14 11.04 & $-$39 50 31.2 & 45000\\\\ \\tableline \\tableline \\end{tabular} \\end{table} ", + "conclusions": "The results of this work may be summarized as follows: 1) Remarkably good fits (a reduced $\\chi$$^2$ of 1.9 for 487 degrees of freedom for all of the data sets in Table 1) to the X-ray spectrum of G347.3-0.5 as observed by the $\\it{ROSAT}$, $\\it{ASCA}$ and $\\it{RXTE}$ satellites have been obtained by using a combination of non-thermal and thermal models (SRCUT and EQUIL) in XSPEC (see Table 2). The X-ray spectrum of G347.3-0.5 at energies higher than 8 keV cannot be adequately fit by simply extending the power-law fits presented by Slane et al. 1999. 2) We have found evidence for modest thermal emission from this SNR. The energy of this thermal component is rather large ($\\it{kT}$ $\\approx$ 1.4 keV) and appears to be associated with the diffuse emission from the SNR rather than the X-ray luminous rims. At an assumed distance of 6 kpc, the corresponding ambient density is $n$$_H$ $\\approx$ 0.02 cm$^{-3}$, consistent with the works of other authors (Slane et al. 1999, Ellison et al. 2001). 3) For the range of break frequencies determined for the different portions of G347.3-0.5 (1.5 through 2.9 $\\times$ 10$^{17}$ Hz), we follow the example of Reynolds \\& Keohane (1999) and calculate a maximum electron energy of $\\approx$ 70 TeV, assuming an ambient magnetic field strength of 10 $\\mu$G. This result is larger than the maximum electron energy calculated with different models for G347.3-0.5 by Ellison et al. 2001." + }, + "0112/astro-ph0112232_arXiv.txt": { + "abstract": "The Wisconsin H-Alpha Mapper (WHAM) has completed a one-degree resolution, velocity-resolved northern sky survey of {H$\\alpha$}~emission from our Galaxy. The unprecedented sensitivity of the instrument and accurate spectral subtraction of atmospheric features allow us to detect Galactic features as faint as 0.1 Rayleighs (EM $\\approx 0.25$ cm$^{-6}$ pc). This survey allows a direct comparison of the ionized and neutral components of the ISM on a global scale for the first time. All-sky maps of {H$\\alpha$}~emission in select velocity bands highlight the rich kinematic structure of the Galaxy's ionized gas. The full set of data from the WHAM survey is now available at {\\tt http://www.astro.wisc.edu/wham/}. One surprising result is that the high latitude sky of both the ionized and neutral components display marked similarity in the location and radial velocity of emitting regions, especially for many of the previously identified intermediate velocity clouds. Although there is evidence for spatial and velocity correlation, in many cases examined so far there is little evidence for a quantitative correlation between the column density of {H {\\sc I}}~and the emission measure of {H {\\sc II}}. WHAM is also capable of studying the ISM through optical emission lines other than H$\\alpha$. Two directions toward the Perseus arm have been studied in detail through several other optical emission lines. The multiple velocity component structure toward these directions provides a selection of ionized environments for study and shows interesting variations in the ratios of these lines. WHAM also has the ability to study select regions of the sky at high spatial resolution (3$'$ to 5$'$) with high velocity resolution. Such observations will allow an even closer comparison between the neutral and ionized components with the advent of high resolution {H {\\sc I}}~surveys. ", + "introduction": "The WHAM instrument is a fully remotely operated facility with a 15 cm, dual-etalon Fabry-Perot spectrometer at the focal plane of a 0.6 m telescope atop Kitt Peak in Arizona (Reynolds et al. 1998). The WHAM spectrometer has a 1$^\\circ$ diameter circular field of view on the sky, and a velocity resolution of 12 km\\thinspace s$^{-1}$~within a 200 km\\thinspace s$^{-1}$~wide spectral window that can be centered on any wavelength between 4800 \\AA~and 7300 \\AA. WHAM was designed to detect very weak emission lines from ionized gas. The Wisconsin H-Alpha Mapper Northern Sky Survey (WHAM-NSS) is the first deep,velocity-resolved survey of interstellar {H$\\alpha$}~emission over the northern sky ($\\delta \\ge -30\\deg$). The survey consists of 37,565 individual observations taken over a span of two years. Each observation recorded the composite spectrum of a one-degree diameter patch on the sky. The spectral resolution of 12 km\\thinspace s$^{-1}$~made it possible to separate cleanly the interstellar emission from the terrestrial emission in every spectrum and to measure the thermal and non-thermal motions of the interstellar gas. Details about the WHAM-NSS and downloadable versions of the survey can be found at {\\tt http://www.astro.wisc.edu/wham/}. These survey maps of interstellar {H$\\alpha$}~emission provide the first global view of the distribution and motions of wide spread ionized hydrogen within the Milky Way (see Figure 1). The small bright knots of {H$\\alpha$}~are classical emission nebulae in the vicinities of hot O and B stars located mostly near the Galactic midplane. Between these bright knots and filling most of the sky is fainter {H$\\alpha$}~emission from the Warm Ionized Medium (WIM), with a characteristic temperature T$_e \\approx~$10$^4$\\thinspace K and density n$_e \\approx~0.1$~cm$^{-3}$. Past studies have shown the WIM to be a significant component of the interstellar medium, especially in the halos of disk galaxies. In the Milky Way, the WIM has a mass surface density about one-third that of neutral atomic hydrogen, a power requirement equal to the kinetic energy injected into the Galaxy by supernovae, and a characteristic scale height above the midplane of 1000 pc, approximately five times larger than that of the neutral hydrogen. This survey of the WIM shows rich structure both on the sky and in velocity, allowing an exploration of the origin of the ionization and heating of this gas and its relationship to the other components of the interstellar medium. ", + "conclusions": "" + }, + "0112/astro-ph0112554_arXiv.txt": { + "abstract": "Please answer ``yes'' or ``no'': \\ \\begin{enumerate} \\item Does the mass function for clusters of galaxies cut off exponentially? \\item Does the luminosity function for galaxies cut off exponentially? \\item Is the dependence of virial velocity on galaxy luminosity a power law? \\item Does the velocity function for galaxies cut off exponentially? \\end{enumerate} ", + "introduction": "The luminosities of cosmic lighthouses are limited by a variety of physical processes. The Eddington limit immediately comes to mind, but like other such limits it is a function of mass. The question of limiting luminosities quickly becomes one of limiting masses. At least four classes of cosmic lighthouses are on the program for this meeting and each has a different typical mass associated with it. It is my non-expert impression that for both stars and AGN we have some understanding of the physics that limits their masses. I will argue that the mass scale for clusters of galaxies is readily explained, but that our understanding (or at least my understanding) of the mass scale for galaxies is very incomplete. ", + "conclusions": "The answers to the four questions posed in the abstract are yes, yes, yes and yes-and-no. The observed velocity function for galaxies appears to cut off exponentially while the velocity function for N-body sub-halos appears not to. There is as yet no detailed physical model which cuts off the condensation of baryons into dark matter halos in a manner which conforms to the observations." + }, + "0112/astro-ph0112281_arXiv.txt": { + "abstract": "We use published mid-IR and V-band flux ratios for images A and B of Q2237+0305 obtained in September 1999 to demonstrate that the size of the mid-IR emission region has a scale comparable to or larger than the microlens Einstein Radius (ER) ($\\sim 10^{17}$cm for microlensing by solar mass stars). Q2237+0305 has been monitored extensively in the R and V-bands for $\\sim$15 years. The variability record shows significant microlensing variability of the optical emission region, and has been used by several studies to demonstrate that the optical emission region is much smaller than the ER for solar-mass objects. For the majority of the monitoring history, the optical flux ratios have differed significantly from those predicted by macro-models. In contrast, recent observations in mid-IR show flux ratios similar to those measured in the radio, and to predictions of some lens models, implying that the mid-IR flux is emitted from a region that is at least 2 orders of magnitude larger than the optical emission region. We have calculated the likeli-hood of the observed mid-IR flux ratio as a function of mid-IR source size given the observed V-band flux ratio. The expected flux ratio for a source having dimensions of $\\sim$1 ER is a sensitive function of the macro model adopted. However we find that the probability of source size given the observed flux ratios is primarily sensitive to the ratio of the macro-model magnifications. Limits on the mid-IR source size can therefore be considered as a function of a one dimensional, rather than a 4 dimensional (two optical depths plus two shears) class of models. The majority of published macro models for Q2237+0305 yield a flux ratio for images B and A of 0.8 - 1.1. By combining probabilities from the ratios A/B and C/D we infer that the diameter of a circular IR emission region is $>1$ER with $>95\\%$ confidence. For sources of this size, other geometries, specifically an annular geometry, appropriate for a dusty torus, yield the same limit if the projected area rather than radius is considered. For microlensing by low-mass stars, this source size limit rules out non-thermal processes such as synchrotron as mechanisms for mid-IR emission. ", + "introduction": "Q2237+0305 (The Einstein Cross) was discovered in the CfA Redshift survey (Huchra et al. 1985). The object comprises a source quasar with a redshift of $z=1.695$ that is gravitationally lensed by a foreground galaxy ($z=0.0394$) producing 4 images with separations of $\\sim 1''$. As a result of the proximity of the lensing galaxy, Q2237+0305 provides a unique opportunity to contrast dynamical measurements based on the geometry of the lensing with more traditional techniques (e.g. Rix, Schneider \\& Bahcall 1992, hereafter RSB92). In addition, the close proximity results in a short time-scale for microlensing, and a large projected microlensing length-scale (with respect to the source). Each of the 4 images are observed through the bulge of a galaxy which has an optical depth in stars that is of order unity (e.g. Kent \\& Falco 1988, hereafter KF88; Schneider et al. 1988, hereafter S88; Schmidt, Webster \\& Lewis 1998, hereafter SWL98). This results in a high probability for microlensing. Microlensing is indicated either by independent temporal variability of image fluxes, or by variation in colour between images at a single epoch (separated by the macro-image delay). Large, and rapid variation of the continuum flux is found in the variability record for Q2237+0305 (Irwin et al. 1989; Corrigan et al. 1991; $\\O$stensen et al. 1996; Wozniak et al. 2000a,b). This variation has been used to argue that the optical emission region must be significantly smaller than the microlens Einstein Radius (ER), and therefore the typical scale of the caustic structure (e.g. Wambsganss, Paczynski \\& Schneider 1990; Wyithe, Webster, Turner \\& Mortlock 2000; Yonehara 2001). During a caustic crossing, a small source exhibits colour variability if the emission spectrum is scale dependent (e.g. Wambsganss \\& Paczynski 1991; Fluke \\& Webster 1999). Evidence of this effect from broad band observations of a microlensing event was presented by Corrigan et al. (1991). Furthermore, if the quasar emits in one wave-band at a scale much smaller than the ER, and on a scale larger than the ER in another band, then colour change may be seen in two random observations of a single image (particularly if the observations straddle a caustic crossing event), or between two different images, as a result of magnification of the smaller source. Ground based observations have confirmed differential amplification of the emission region. Lewis et al. (1998) determined the ratios of emission line equivalent widths relative to one image. They show $(i)$ the ratios remain fairly constant for one image from line to line, suggesting that the sizes of the emission regions for the lines are not greatly different, $(ii)$ that the ratios vary from image to image for a single epoch by a factor of $\\sim 2.5$, and $(iii)$ that the ratio for a single image varies as a function of time, i.e. as a result of a microlensing event. These results are consistent with earlier results of Fillipenko (1989) who measured a $\\sim25\\%$ difference in the width of the MgII lines between the A and B images. The CIII] line, produced by extended broad line regions has been measured in an attempt to find the flux ratios $(R)$ using emission scales beyond the influence of microlensing (Yee \\& De Robertis 1992; Racine 1992; Fitte \\& Adam 1994; Saust 1994; Lewis et al. 1998). Mediavilla et al. (1998) observed the CIII] line in Q2237+0305 using two-dimensional spectroscopy, and found an arc (image of the extended narrow line region of the source) extending around three of the images, indicating a very extended region of emission. On the other hand, measurement of $R$ using the CIII] line is subject to differential extinction, and uncertainties in continuum subtraction. To avoid the effects of extinction, and possibly microlensing Falco et al. (1996) imaged Q2237+0305 at 3.6cm, finding flux ratios similar to those inferred from CIII]. In addition, Q2237+0305 has also been imaged in the Ultraviolet (Blanton, Turner \\& Wambsganss 1998), and in X-rays (Wambsganss, Brunner, Schindler \\& Falco 1999). Many models have been proposed for the projected lens mass distribution based on observations of the lensed images of Q2237+0305 (e.g. KF88; S88; Kochanek 1991, hereafter K91; Wambsganss \\& Paczynski 1994, hereafter WP94; Witt, Mao \\& Schechter 1995, hereafter WMS95; SWL98; Chae, Turnshek \\& Khersonsky 1998, hereafter CTK98). The majority of these predict a flux ratio for images B and A of $R_{BA}\\sim 0.80 - 1.1$, consistent with the ratio measured in the radio of $R_{BA}=1.1\\pm0.3$ (Falco et al. 1996). However, over the monitoring history, the optical light-curve shows variations in $R_{BA}$ between $\\sim 0.2$ and $\\sim 1.0$. The discrepancy is attributed primarily to microlensing, and firmly demonstrates that the optical flux ratios cannot be used as model constraints (e.g. S88; KF88; K91; WP94). Agol, Jones \\& Blaes (2000) (hereafter AJB00) have found a mid-IR B:A flux ratio of $R_{BA}\\sim1.1$. This ratio is consistent with observations in the radio (Falco et al. 1996). AJB00 interpreted their results as evidence for an extended region of mid-IR emission, with dimensions larger than the microlens Einstein Radius. In this paper we use microlensing models to calculate distributions of flux ratios for sources with different sizes and intensity profiles, and hence derive quantitative limits on the scale of the mid-IR emission. In Secs.~\\ref{models} and \\ref{macros} we describe the microlensing models, and summarise published macro-models for Q2237+0305. Sec.~\\ref{results} discusses the methods used to infer the mid-IR source size from the observed optical and mid-IR flux ratios, and the source size limits implied by the published macromodels. Initially we restrict our discussion to the particular case of the flux ratio between images B and A. However we present results based on all image ratios in Sec.~\\ref{sec_pairs}. In the conclusion we mention some implications for quasar physics. ", + "conclusions": "Monitoring of the gravitationally microlensed Quasar Q2237+0305 has found significant microlensed optical variability. Observations in July and September of 1999 showed (de-reddened) V-band flux ratios between the different sets of images that differed by large factors from the corresponding flux ratios in the mid-IR. In addition, the mid-IR flux ratios are similar to those measured in the radio, and to predictions of some lensing models, suggesting that mid-IR emission region is not subject to large microlensing variation. The mid-IR emission region is therefore larger than the microlens Einstein Radius ($\\eta_o$), and hence at least 2 orders of magnitude larger than the optical emission region, which is thought to be $<0.01\\eta_o$ (Wambsganss, Paczynski \\& Schneider 1990; Wyithe, Webster \\& Turner 2000). The colour difference between different images is due to the magnification/de-magnification of the optical emission. We have used microlensing models to calculate, as a function of mid-IR source size, the probability of obtaining the observed mid-IR flux ratios given the observed optical flux ratios. An alternative approach might be to note that the V-band flux in some images varied between the two mid-IR observations, while the mid-IR flux remained steady. In this case, the optical ratios need not be de-reddened, and as a result carry smaller errors. However, the variation in V-band is not significantly larger than the observational uncertainty in the mid-IR. We have therefore restricted our attention to flux ratio variation between different images. The flux ratio statistics for large sources are sensitive to the macro-model assumed. We have therefore computed source size limits for an ensemble of models. The probability has the interesting property of being primarily sensitive to the macro-model flux ratio. As a result we have parameterised the macro-models available from the literature by their predicted flux-ratio. The strongest limits are found from the flux ratios between images B and A. This is due in part to the differences between the relative V-band and mid-IR fluxes, as well as the more consistent macro-model predictions for these images. We find that the mid-IR source size $S_{IR}>1\\eta_o$ with $>90\\%$ confidence, and $>0.5\\eta_o$ with $>95\\%$ confidence. The IR-emission scale is larger than the optical emission with a confidence $>99\\%$. The limit on the infrared source size derived here for the Einstein Cross may be converted to a limit on the brightness temperature. Assuming microlensing by stars, $S_{IR} \\ga \\eta_o \\sim 10^{17}$cm, and the brightness temperature at 10 $\\mu$m (rest frame 3.7 $\\mu$m) is about $T_b \\la 7900 K$ for a luminosity distance of $10^{28} $cm, magnification of 15, and flux of 20 mJy. This upper limit on the source brightness rules out non-thermal emission mechanisms, such as synchrotron, which typically have brightness temperatures of $10^{8-10}$K. As argued in AJB00, the spectrum indicates thermal emission at $\\sim 2000$K, which is comparable to the sublimation temperature of dust. The flux from the QSO is sufficient to heat the dust at the sublimation radius of $\\sim$~1~pc. One can turn this argument around: assuming that the IR emission is due to dust at the sublimation radius, then one can estimate the variance of the flux of each image for the given source size and macro lens model. We can then compute the $\\chi^2$ of each macrolensing model using this variance. Analytic estimates of the variance for large source exist for models with $\\gamma=0$ \\cite{R93} and some numerical estimates exist for non-zero $\\gamma$ \\cite{R97}. The estimates with zero $\\gamma$ indicate that the variance should be of order 10\\% or less, unfortunately about the same size as the observational error bars for the infrared observations. Since models for a large source size require much larger ray-tracing simulations, with a lower spatial resolution than we have carried out, we leave this computation for future work. Future observations with higher signal-to-noise, or a measurement of the variance with long-term IR monitoring will allow a better estimate of which macro-lens model is correct." + }, + "0112/astro-ph0112248_arXiv.txt": { + "abstract": "We have carried out an in-depth study of the peripheral region of the molecular cloud L1204/S140, where the \\fuv\\ radiation and the density are relatively low. Our observations test theories of photon-dominated regions (PDRs) in a regime that has been little explored. Knowledge of such regions will also help to test theories of photoionization-regulated star formation. \\CII\\ 158 \\micron\\ and \\OI\\ 63 \\micron\\ lines are detected by ISO at all 16 positions along a 1-dimensional cut in right ascension. Emission from \\hh\\ rotational transitions \\jj20\\ and \\jj31, at 28 and 17 \\micron, was also detected at several positions. The \\CII, \\OI, and \\hh\\ intensities along the cut show much less spatial variation than do the rotational lines of \\CO\\ and other CO isotopes. The average \\CII\\ and \\OI\\ intensities and their ratio are consistent with models of PDRs with low \\fuv\\ radiation (\\go) and density. The best-fitting model has $\\go \\sim 15$ and density, $n \\sim \\eten{3}$ \\cmv. Standard PDR models underpredict the intensity in the \\hh\\ rotational lines by up to an order of magnitude. This problem has also been seen in bright PDRs and attributed to factors, such as geometry and gas-grain drift, that should be much less important in the regime studied here. The fact that we see the same problem in our data suggests that more fundamental solutions, such as higher \\hh\\ formation rates, are needed. Also, in this regime of low density and small line width, the \\OI\\ line is sensitive to the radiative transfer and geometry. Using the ionization structure of the models, a quantitative analysis of timescales for ambipolar diffusion in the peripheral regions of the S140 cloud is consistent with a theory of photoionization-regulated star formation. Observations of \\CII\\ in other galaxies differ both from those of high \\go\\ PDRs in our galaxy and from the low \\go\\ regions we have studied. The extragalactic results are not easily reproduced with mixtures of high and low \\go\\ regions. ", + "introduction": "Photon-dominated regions (or photodissociation regions, PDRs) are regions of the neutral interstellar medium (ISM) where far-ultraviolet (FUV) (6 eV $< h\\nu < 13.6$ eV) photons control the heating and chemical processes. They are the interface between \\HII\\ regions and cold molecular cores. The physical and chemical structure of PDRs depends critically on the FUV intensity and the gas density. In the peripheral regions of the molecular cloud L1204/S140, the FUV intensity and the gas density are low, allowing tests of the models in an important regime. We are motivated by two primary goals: understanding PDRs in a regime of parameter space that has not been extensively studied; and understanding the role of the regions with relatively low FUV and density in global star formation, both in our own Galaxy and in other galaxies. A great deal of observational and theoretical effort has been devoted to understanding PDRs. Comprehensive models of PDRs have been constructed by several groups (e.g., Black \\& Dalgarno 1977, van Dishoeck \\& Black 1986, 1988, Tielens \\& Hollenbach 1985a, Sternberg \\& Dalgarno 1989, le Bourlot et al.\\ 1993, Kaufman et al. 1999) by solving the full chemistry and heating-cooling balance in a self-consistent way. In PDRs with $n \\leq 10^5$ \\cmv, the most important heating process for gas is photoelectric heating---electrons ejected from dust particles by FUV photons heat the gas through collisions. For dust, direct absorption of FUV photons is the primary heating mechanism. Far-infrared continuum emission is the major cooling process for dust, and the \\fir\\ lines of \\CII\\ $^2$P$_{3/2}\\rightarrow ^2$P$_{1/2}$ at 157.7409 \\micron\\ (hereafter \\CII) and \\OI\\ $^3$P$_1\\rightarrow ^3$P$_2$ at 63.183705 \\micron\\ (hereafter \\OI) are the most important for gas. Therefore, the \\CII\\ and \\OI\\ lines, along with \\hh\\ rotational emission, are the most important tracers of PDRs. Most previous work has focused on bright PDRs very close to hot OB stars, e.g., the Orion bar (Tielens \\& Hollenbach 1985b; Jansen et al.\\ 1995; Hogerheijde et al.\\ 1995; Tauber et al.\\ 1994), the NGC 2023 PDR (Steiman-Cameron et al. 1997, Draine \\& Bertoldi 1996), and the S140 PDR (Emery et al. 1996; Timmermann et al.\\ 1996; Spaans \\& van Dishoeck 1997). Other recent studies include those by Liseau et al. (1999), and the field has been reviewed by Draine \\& Bertoldi (1999) and by Hollenbach \\& Tielens (1999). These regions have a \\fuv\\ intensity of \\go\\ $> 10^3$ and a density higher than 10$^4$ \\cmv, where \\go\\ is the enhancement factor relative to the standard interstellar radiation field as given by Habing (1968). There has been very little exploration of the physics of PDRs with modest \\fuv\\ fields and densities, conditions likely to prevail over most of the surface of molecular clouds in our Galaxy. Federman et al. (1995) and van Dishoeck \\& Black (1988) have studied PDRs in diffuse clouds ($A_V \\sim 1$ mag) and translucent clouds ( $n< 1000$ \\cmv, \\go\\ $< 17$ and $A_V < 5$ mag). Regions with high densities and moderate UV fields ( $n< 5\\ee4$ \\cmv, $\\go\\ = 900$) have also been studied in some detail (Jansen et al. 1995), and Kemper et al. (1999) have used \\submm\\ and \\fir\\ observations to probe a reflection nebula with $n \\sim 5000$ \\cmv\\ and $ \\go \\sim 200$. In this paper, we explore the critical intermediate regime where $n \\equiv\\ n({\\rm H}) + 2 n({\\rm H_2}) \\sim 500-5000$ \\cmv\\ and $\\go \\sim 10-60$. The Infrared Space Observatory (ISO\\footnote{ISO is an ESA project with instruments funded by ESA Member States (especially the PI countries: France, Germany, the Netherlands and the United Kingdom), with the participation of ISAS and NASA.}) provided a unique opportunity to observe low-brightness extended \\CII, \\OI, and \\hh. We used this capability to study the intermediate regime. It is also important to understand the role of regions with modest extinction and column density in star formation. Regions like those we are studying include most of the mass in the interstellar medium (Hollenbach \\& Tielens 1995), but their role in star formation is poorly known. Based on the Jeans criterion, most molecular clouds in the Galaxy are not sufficiently supported by thermal energy and therefore should collapse under gravity to form stars. Such widespread collapse, however, would lead to a Galactic star formation rate hundreds of times higher than what has been observed (Zuckerman \\& Palmer 1974, Evans 1991). The observed low star formation rate seems to indicate that most parts of most molecular clouds are ``sterile'' (Evans 1999). Magnetic fields and turbulence are generally considered to play an important role in supporting molecular clouds and preventing or slowing collapse. However, there is no widely accepted theory on the mechanism of magnetic and turbulent support of molecular clouds. Recently, Elmegreen (2000) has argued that star formation does in fact happen within a few cloud crossing times, removing the need for cloud support. Pringle, Allen, \\& Lubow (2001) have pointed out that such a picture strains methods of cloud formation, and they conclude that visible clouds would have to form out of ``dark\" molecular matter. These scenarios beg the question: what prevents star formation in the great majority ($\\sim 98$\\%) of molecular gas? McKee (1989; Bertoldi \\& McKee 1996, 1997) proposed a mechanism of photoionization-regulated star formation to explain the low star formation rate in the Galaxy by combining two ideas: first, that magnetic fields support molecular clouds; and second, that PDRs occupy a large fraction of molecular gas. The timescale for ambipolar diffusion is proportional to the ionization fraction ($\\xe = n(e)/n$) and the FUV photons and cosmic rays are the sources of ionization. Most molecular gas is sterile because it resides in a PDR, where the ionization is high enough to slow ambipolar diffusion. Only highly shielded regions are likely to form stars. In addition, newly formed stars inject energy into the clouds, replenishing turbulence and slowing star formation in the rest of the cloud. In this picture, molecular clouds reach dynamic equilibrium when $\\AV \\sim 8$ mag. By focusing on a peripheral region, we test the conditions in these regions, which should not be forming stars according to the theory. We have chosen a 1-D positional cut---we call it the S140 cut---in the peripheral region of the molecular cloud L1204 (Figure~\\ref{s140-iso}) for our PDR and chemistry study. S140, at distance of $\\sim\\! 910$ pc (Crampton \\& Fisher 1974), is an \\HII\\ region associated with the molecular cloud L1204. In the rest of the paper, we will use S140 to refer to the whole cloud. The B0V star HD 211880 illuminates this cloud from the southwest side to create the visible \\HII\\ region, an ionization front, and a bright PDR that has been studied extensively. For a picture that shows the CO emission in relation to the ionization front and exciting star, see Fig. 1 of Zhou et al. (1994). Northeast of the bright PDR lies the most prominent dense core (hereafter referred to as the S140 dense core) associated with the cloud. Numerous studies have been carried out on this cloud, e.g., Blair et al. (1978), Tafalla, Bachiller, \\& Martin-P\\'{\\i}ntado (1993), Plume, Jaffe, \\& Keene (1994), Emery et al. (1996), and Timmermann et al. (1996). The FUV radiation at the ionization front provided by HD 211880 is about 150 times Habing's mean interstellar radiation field (Keene et al.\\ 1985; Habing 1968). The S140 cut, specified in Table~\\ref{isopos}, is far enough ($\\sim 20 \\arcmin$) from HD211880 to have a low FUV radiation field ($\\go \\sim 15 - 60$, \\S \\ref{sec4:fuv}, \\ref{sec4:kaufman}) and modest visual extinction ($A_V \\sim 15$, \\S \\ref{sec4:density}). ", + "conclusions": "\\label{sec4:model} In this section, we will attempt to reproduce the observations by using both published models and our own PDR models. The most important free parameters in the modeling are the FUV radiation intensity, the total extinction, and the volume density. Therefore, we will first constrain the FUV intensity (\\S~\\ref{sec4:fuv}) and the extinction and density (\\S~\\ref{sec4:density}). Then we will compare the predicted \\CII\\ and \\OI\\ line intensities with the ISO observations in \\S~\\ref{sec4:kaufman}, \\ref{hst}. \\subsection{The FUV Radiation Field Estimated from the Far-infrared Continuum} \\label{sec4:fuv} Because dust grains absorb most of the FUV photons incident on the cloud and most of that energy is re-emitted in the form of \\fir\\ continuum, the intensity of the FUV radiation can be inferred from the \\fir\\ continuum. Hollenbach et al.\\ (1991) presented model predictions of the \\fir\\ continuum intensities at 100 \\micron\\ and 60 \\micron\\ for a range of incident FUV intensity. {}From Table~\\ref{s140lws-res} and Figure~\\ref{s140cut-lines}, the IRAS 100 \\micron\\ and 60 \\micron\\ intensities are quite uniform along most of the S140 cut except at Positions 15--16 for 60 \\micron\\ and the Positions 13--16 for 100 \\micron, where the intensities are lower. Therefore we use the intensities averaged over positions 1-13 and the intensity at position 16 to estimate a range of values for the ultraviolet intensity. The continuum intensities at the ISO Off position have been subtracted to yield $I(60) = 30 - 50$ MJy sr$^{-1}$ and $I(100) = 100 - 180$ MJy sr$^{-1}$. Assuming a $\\lambda^{-1}$ dust emissivity, the total \\fir\\ intensity is about $6\\times 10^{-3}$ ergs s$^{-1}$ cm$^{-2}$ sr$^{-1}$. These intensities and a $\\lambda^{-1}$ dust emissivity yield a dust temperature of $(28\\pm 2)$ K, near the constant value of 27 K reached for $\\go < 160$, when effects of transient heating of small grains/PAHs by single photons are included (Hollenbach \\& Tielens \\ 1995). Therefore, the 60 \\micron\\ continuum intensity in this region is probably strongly affected by small grains. Based on Fig. 18 of Hollenbach et al. (1991), we used the 100 \\micron\\ intensity to estimate the strength of the ultraviolet field as $\\go = 40 - 60$. As noted by Kaufman et al. (1999), these values assume that the FUV impinges only on one surface of the cloud. If the cloud is heated from both sides, then the values of \\go\\ appropriate to one surface would be 20 -- 30. The assumption here is that the cloud is optically thin to the \\fir\\ continuum. Given the uncertainties in the fraction of the grain heating caused by FUV photons, we estimate \\go\\ in the range, 15--60. The enhancement factor of the radiation field $I_{UV}$ with respect to the standard field given by Draine (1978) is 1.71 times smaller than \\go, $I_{UV} = \\go/1.71$ (Draine \\& Bertoldi 1996). \\subsection{Column Density and Density Regime} \\label{sec4:density} The mean integrated intensity of \\COOTWO\\ for the cloud component is 0.54 K \\kms. From Table~\\ref{s140cso-res}, the typical excitation temperature of \\CO\\ is about 10 K. Assuming that the \\COOTWO\\ line is optically thin and in LTE, we calculate analytically the total column density of \\COO\\ to be $\\leq 1.3\\times 10^{14}$ cm$^{-2}$. The empirical relations of Frerking, Langer, \\& Wilson (1982) on $\\rho$ Oph yield a visual extinction, $A_V = 3.4$ mag, and a total column density of H nuclei of $5\\times 10^{21}$ cm$^{-2}$. Assuming the cloud is spherical and the angular size of the cloud along the S140 cut is roughly 15\\arcmin, the linear size of the depth of the cloud would be $1.2\\times 10^{19}$ cm (4 pc) at a distance of 910 pc. Therefore, the average density along the S140 cut would be 800 \\cmv. The assumption of LTE is not valid at such a low density. Models of excitation, including trapping with LVG codes, indicate that the \\COO\\ populations at such low densities are far from LTE. A more self-consistent solution is a density around 2000 \\cmv\\ and a column density of \\COO\\ of 2.0\\ee{15} \\cma, yielding $A_V = 16$. The ratio of $J=2-1$ to $J=3-2$ lines of \\CCO\\ is roughly consistent with a density of 2000 \\cmv, and the non-detection of \\HCOP\\ $J = 3-2$ outside core F limits the density to $n < 10^4$ \\cmv\\ for typical abundances in translucent clouds and $\\tk = 10-20$ K. Similar considerations applied to the position of core F (774,0) lead to estimates of $A_V = 25 $ mag and, using a size of 3\\arcmin, a mean density of \\eten4 \\cmv. Considering the uncertainties in these estimates, the likely uncertainties in column density and density are factors of 2 and 3, respectively. A direct comparison between PDR models and our data will be given in \\S \\ref{hst}, which supports these simple estimates. \\subsection{Comparison to Published PDR Models} \\label{sec4:kaufman} We first compare our results to the published grid of PDR models by Kaufman et al. (1999). These calculations incorporated new collision rates for fine-structure lines and \\hh, new PAH heating and chemistry, and lower gas-phase abundances for oxygen and carbon (Savage \\& Sembach 1996). They include regions of low FUV and low density, which seem to be appropriate for our situation. The calculations employed a turbulent broadening of 1.5 \\kms, about twice what we infer, but Kaufman et al. argue that the results are not too sensitive to this parameter. We first consider the constraints from the two atomic species, \\CII\\ and \\OI. From Fig. 4 of Kaufman et al., the mean ratio of \\OI/\\CII\\ of $0.16\\pm 0.4$ constrains $\\go < 20$. This is at the low end of the range inferred in \\S \\ref{sec4:fuv}, even if we assume the cloud is heated from both sides. At this \\go, the only solution for density is $n \\sim 300$, lower than the values inferred in \\S \\ref{sec4:density}. For lower \\go, the solutions for $n$ bifurcate into lower and higher values. For example, at $\\go = 10$, there are solutions for log $n = 2.0$ or 3.3. The latter are more consistent with the LVG modeling of \\COOTWO. The \\CII\\ emission by itself (Fig. 3 of Kaufman et al.), is less diagnostic, but roughly consistent with the ranges implied by the ratio. Are these low values of \\go\\ consistent with the \\fir\\ emission? They are substantially below the estimates from \\S \\ref{sec4:fuv}. Our ratio of intensities of the sum of the \\OI\\ and \\CII\\ lines to the total \\fir\\ intensity, 1.7\\ee{-2}, is close to the highest values found by Kaufman et al. (1999). In fact, the ratio of line to continuum of the cloud in S140 may exceed the maximum in their models, if indeed the cloud is heated from both sides. What about the \\CI\\ data? The mean for the cloud, summing the two components as for the \\COO\\ data, and correcting for efficiency, is 1.3\\ee{-6} \\icgs. Comparing to Fig. 7 of Kaufman et al., this value favors very low \\go. To get \\go\\ up to even 15 requires $n \\sim 100$. For the conditions that match the other lines, the \\CI\\ prediction of Kaufman et al. is too high by a factor of 3--4. Given reasonable uncertainties, the mean cloud data require quite low values of \\go\\ and $n$, with most likely values of $\\go \\sim 15$ and $n \\sim 1000$. These are both lower than our initial guesses, even if the cloud is heated from both sides. For this range of conditions, the ``surface temperature\" of the PDR, plotted in Fig. 1 of Kaufman et al., will lie between 100 and 200 K. This is the maximum temperature that should apply to the region of \\hh\\ emission, if collisions dominate the \\hh\\ excitation. The excitation temperatures of \\hh\\ (Table \\ref{s140sws-res}) are consistent with this range of temperatures, but the absolute intensities are difficult to reproduce (\\S \\ref{sec4:h2interp}). Comparison of the Kaufman et al. results with those of other models in the literature (e.g., Le Bourlot et al.\\ 1993, Roueff, private communication) and our own models (e.g., Jansen et al.\\ 1995, Spaans \\& van Dishoeck 1997, Spaans, private communication), shows good agreement for the \\CII\\ intensities but variations of factors of two in the \\OI\\ intensities depending on the temperature structure, radiative transfer treatment and geometry of the source. Thus, the range in inferred \\go\\ and $n$ could be somewhat larger, but the conclusion that both are low holds firm. The region of the core F emission is clearly denser than that of the cloud, yet core F is not apparent in either \\fir\\ lines or continuum. The emission from the \\CII\\ line may be quite independent of density up to about $n \\sim \\eten4$ \\cmv, in the allowed range of \\go; however, the \\OI\\ line should increase with density. From Fig. 4 of Kaufman et al. (1999), the ratio of \\OI/\\CII\\ should increase to about 1 for $n = \\eten4$ and \\go\\ in the range that fits our data. To keep the ratio below 0.3 would require $\\go \\sim 1$. That is, core F would have to be shielded from the ambient FUV field. Interestingly, the \\CI\\ emission does see core F in the $-10.5$ \\kms\\ line. Either that component is not completely shielded or the enhanced \\CI\\ is caused by some internal source that does not affect the \\fir\\ line and continuum emission. \\subsection{Comparsion to PDR-Monte Carlo Models} \\label{hst} To check the effects of radiative transport, different Doppler broadening, and metal abundance at these low values of \\go\\ and $n$, we ran our own models for the part of parameter space indicated in the last section. We used our own PDR code, an updated version of the code described by Jansen et al. (1995) to calculate abundances of relevant species as a function of depth into the cloud. We then fed the results into a Monte Carlo code (Hogerheijde \\& van der Tak 2000) to calculate excitation and radiative transport. Einstein A values of \\CII\\ and \\OI\\ were taken from Tielens \\& Hollenbach (1985a) and Galavis et al. (1997), respectively. In going from the PDR code to a spherical Monte Carlo code, one doubles the column density, so the PDR models were run for $\\AV = 9$, yielding a total extinction through the cloud of 18 mag. We ran models for Doppler parameters from 0.5 to 2.0. The integrated intensities were generally insensitive to this parameter. Two sets of values for metal abundances were used: normal abundances are those listed by Jansen et al. (1995) in their Table 2 for the $\\zeta$ Oph diffuse cloud; ``low metals\" corresponds to decreasing all abundances other than C, N, and O by a factor of 10. The most important species is S, which is thought to be quite depleted in molecular clouds compared to translucent clouds. The results showed that the radiative transport of \\OI\\ is indeed quite sensitive to the input parameters because it is very subthermally excited at these low densities, but quite opaque ($\\tau \\sim 10$ for models that fit the data well). At very low values of \\go\\ and $n$, the \\OI\\ line predicted by the Monte Carlo cloud can be considerably stronger than that predicted by the PDR code alone, after accounting for the doubling of the column density. The differences seem to be partially in the handling of the radiative transport and partly in the geometry. Calculations that calculate the temperature self consistently in spherical geometry produce even larger effects (M. Spaans, personal communication). Among a grid of models, the best fit was obtained for a model with $\\go = 17$ and $n = 1000$ and low metals. This model reproduced the \\COO\\ line strength, indicating that the extinction estimate from \\S \\ref{sec4:density} was about right. The ratio of \\CCOTHR/\\CCOTWO\\ was about 50\\% lower than the observations in this model, while models with $n = 3000$ gave ratios higher than observed and \\CCOTWO\\ lines about a factor of 2 stronger than observed. The \\HCOPTHR\\ line produced by this model is about 1.5 times our RMS noise, still consistent with the observations, while the model with $n = 3000$ would produce a 10 $\\sigma$ detection. The \\HCOP\\ abundance and thus the \\HCOPTHR\\ line are considerably enhanced in the case of low metals. Taken together, these results suggest that $n = 1000$ \\cmv\\ is about right, but perhaps a bit low. The best model gave a \\CII\\ line in good agreement with observations and an \\OI\\ line about 40\\% too strong. Higher values of \\go\\ or $n$ greatly overproduced the \\OI\\ line. These are in rough agreement with Kaufman et al., but the \\CI\\ line predicted by our models is about a factor of 4 weaker than the line predicted by Kaufman et al., and our prediction agrees with the observations. Kaufman et al. note that their \\CI\\ lines are stronger than in previous calculations because of inclusion of PAH chemistry in their models. While charge transfer from PAHs to \\CP\\ to create \\CZ\\ is included in our models, the details may be different. We use the results of the model with $\\go = 17$ and $n = 1000$ in the next section. \\label{sec4:sum} We have studied the peripheral region of the molecular cloud L1204/S140, where the FUV radiation and density are relatively low, using ISO to observe \\CII, \\OI, and \\hh\\ lines, and the CSO to observe the submillimeter line of atomic carbon and millimeter emission lines of \\CO\\ and other CO isotopes. We analyzed the results with published PDR models and our own PDR models, coupled to a Monte Carlo simulation code for radiative transfer. The important results from this work are: \\begin{enumerate} \\item We detected wide-spread \\CII, \\OI, and \\hh\\ emission in the region. The \\CII\\ and \\OI\\ lines display much less spatial variation than \\CI\\ and the lines from \\CO\\ and other CO isotopes, which reveal the presence of a denser, shielded core that does not emit significantly in the \\CII\\ and \\OI\\ lines. \\item The average \\CII\\ and \\OI\\ intensities of the S140 cut are $(8.66\\pm 1.27)\\times 10^{-5}$ and $(1.35\\pm 0.34)\\times 10^{-5}$ ergs s$^{-1}$ cm$^{-2}$ sr$^{-1}$; the average line ratio (\\OI/\\CII) is $0.16\\pm0.04$. Therefore the \\CII\\ line is the dominant cooling line in this region. This is quite different from what has been observed in the S140 bright PDR by ISO, where the ratio \\OI/\\CII\\ is $1.55\\pm1.10$. \\item Using the model predictions in Hollenbach et al.\\ (1991) and the 100 \\micron\\ IRAS flux, the FUV intensity \\go\\ for the region would be 40--60, in the situation of single-sided FUV radiation on the cloud. If the FUV impinges on both sides, the inferred $\\go \\sim 20-30$. Based on \\COO\\ emission and simple LVG modeling, the density of the gas, $n \\sim 1000$ \\cmv, and $\\AV \\sim 16$. Conditions in the shielded core are different: $n \\sim 1\\ee4$ \\cmv, and $\\AV \\sim 25$. \\item The PDR models (both published and our own) indicate that the observed \\OI\\ intensity and the ratio \\CII/\\OI\\ are best reproduced with $\\go \\sim 15$ and $n \\sim 1000$ \\cmv. Our models, coupled to a Monte Carlo simulation of the radiative transfer, reproduce the observed \\COOTWO\\ emission for $\\AV \\sim 18$ mag. These conditions are also consistent with the constraints on density from line ratios of \\CCO\\ and \\HCOP. Our models also reproduce the observed \\CI\\ line, while the published models of Kaufman et al. (1999) overproduce this line. In this regime of $n$ and $\\go$, the \\OI\\ intensity is very difficult to model because it is extremely subthermal and also very opaque (the optical depth is 9--14). Trapping is very important in determining the emergent \\OI\\ intensity. \\item Quantitative analysis of the ionization structure and timescale for cloud collapse for the peripheral regions of S140 leads to conclusions consistent with a picture of photoionization-regulated star formation. \\item The detection of widespead \\hh\\ rotational emission is not readily explained with standard PDR models and \\hh\\ formation rates. Possible solutions include a substantial enhancement in the the \\hh\\ formation rates or localized heating by weak shocks associated with turbulent decay. \\item While the emission in \\CII\\ and \\fir\\ continuum from other galaxies lies between the characteristics of the emission from the peripheral region studied here and regions of higher \\go\\ and $n$, simple mixtures of these regions are not able to reproduce the characteristics of the emission from other galaxies. \\end{enumerate} We thank the referee for a careful reading and for suggestions that improved the paper. We would like to extend our sincere thanks to the wonderful service by the ISO supporting staff at IPAC (S. Lord, S. Unger, D. Levine, L. Hermans especially) and at SRON Groningen (E. Valentijn and F. Lahuis in particular). We also thank the CSO staff members and E. Gregersen and K. Mochizuki for helping with the CSO observations. M. Spaans and D. Jansen provided useful discussions and help with PDR codes; E. Roueff provided data for comparison of PDR models. This research has made use of NASA's Astrophysics Data System Abstract Service, the Simbad database, operated at CDS, Strasbourg, France, and the Online Services of IRSKY and IBIS at IPAC, and the SkyView Online Service. WL was partially supported by a Continuing Fellowship and a David Bruton, Jr., Fellowship of the University of Texas, a Frank N. Edmonds, Jr. Memorial Fellowship and a David Alan Benfield Memorial Scholarship of the Department of Astronomy, the University of Texas. Work with the CSO re-imager was supported by NSF grant AST-9530695. The research was supported by NASA Grants NAG2-1055, NAG5-3348, and the State of Texas, NWO grant 614.41.003, a NWO Spinoza grant, and a NWO bezoekersbeurs. \\clearpage" + }, + "0112/astro-ph0112412_arXiv.txt": { + "abstract": "The observed magnetic field of microgauss strength in clusters of galaxies should induce the Faraday rotation effect on the linearly polarized cosmic microwave background (CMB) radiation when the CMB radiation propagates through a cluster at low redshift. The Faraday rotation arises from combined contributions of the magnetic field strength, the electron density, the cluster size, and the characteristic scale of the magnetic field along the line of sight. Employing the Press-Schechter prescription for the cluster abundance under the cold dark matter (CDM) scenario and a plausible isothermal $\\beta$-model for the gas distribution, we compute angular power spectra of the CMB polarization fields including the Faraday rotation mixing effect under the simple assumption of uniform magnetic field configuration across a cluster. As a result, we find that a {\\em parity-odd} $B$-type polarization pattern is statistically generated on the observed sky, even when the primary polarization only contains the {\\em parity-even} $E$-mode component, such as in the case of pure scalar perturbations. The generated $B$-type polarization has a peak with amplitude of $\\sim0.1~\\mu{\\rm K}(B_0/0.1~\\mu{\\rm G})(\\nu_0/10~{\\rm GHz})^{-2}$ at angular scales of $l\\approx 1000$ for the currently favored adiabatic $\\Lambda$CDM model. This result also implies that, if the magnetic field has a $0.5~\\mu{\\rm G}$ strength and we observe at lower frequencies such as $\\nu_0\\simlt 5~{\\rm GHz}$, the secondary signal due to the Faraday rotation effect could be comparable to the magnitude of the primary polarization. The frequency dependence of the Faraday rotation can be then used to discriminate the effect from primary and other secondary signals on the CMB polarizations. Our results therefore offer a new empirical opportunity to measure or constrain the intracluster magnetic field in the average sense, combined with measurements of the intracluster gas distribution through the X-ray and SZ data. ", + "introduction": "\\label{intro} Various observations have revealed that clusters of galaxies are pervaded by the strong magnetic field of microgauss strength (e.g., see \\cite{Kronberg94}; \\cite{Carilli} for reviews and references therein). The multi-frequency Faraday-rotation measurements of polarized radio sources inside or behind a cluster have been used to estimate the magnetic field strength, combined with the X-ray data (\\cite{Vallee}; \\cite{Kim}). Recently, Clarke, Kronberg, \\& B\\\"ohringer (2001) have drawn a firm conclusion that an intracluster hot plasma is universally magnetized by $1-10\\mu$G fields under the assumption of $10-100$kpc field coherent scales, using $16$ normal low-$z$ {\\it ROSAT} cluster sample selected to be free of unusual strong radio halos, widespread cooling flow and recent merger activity. However, except for a few cases such as some background rotation measures per square degree in the Coma cluster (\\cite{Kim90}), it is generally difficult to measure the angular profile of the magnetic field strength inside a cluster because of the lack of the number of radio sources available per cluster. The observations of cluster-wide diffuse radio halos have also led to the evidence of the microgauss field, which is believed to be synchrotron radiation by relativistic electrons accelerated in the shock wave. Moreover, in this case hard X-ray emission could be produced by the same relativistic electron population through the inverse Compton scattering of cosmic microwave background (CMB) photons. The combined observation of hard X-ray emission and radio halos is therefore very attractive in the sense that it allows us to directly estimate the magnetic field strength without further restrictive assumptions for the coherent length and the electron distribution. The recent detections of hard X-ray emission have led to an independent estimation of $\\sim0.1\\mu$G fields even in the outer ($\\simgt 1{\\rm~Mpc}$) envelopes of clusters (\\cite{Bagchi}; \\cite{Rephaeli99}; \\cite{Fusco}). However, the origin of the intracluster magnetic fields still rests a mystery in cosmology. The following two scenarios have been usually investigated in the literature. One is based on the idea that the cluster field is related to the fields generated by dynamo mechanism in individual galaxies and subsequent wind-like activity transports and redistributes the magnetic fields in the intracluster medium (e.g. \\cite{Kronberg99} and references therein). The general prediction of galactic dynamo mechanism is that the galactic magnetic field could arise from an exponential amplification of a small {\\em seed} field during a galactic lifetime and energy of the mean magnetic field could grow up to equipartition level with the turbulent energy of fluid (\\cite{Parker}; \\cite{Zeld}; \\cite{Chiba}). It is not still clear, however, that the galactic dynamo theory can explain the detection of microgauss magnetic fields in high-$z$ damped Ly$\\alpha$ absorption systems, which are supposed to be protogalactic clouds (\\cite{Welter}; \\cite{Wolfe}). On the other hand, an alternative scenario is that the intracluster magnetic fields may grow via adiabatic compression of a primordial field frozen with motions of cosmic plasma, where the primordial field is assumed to be produced somehow during the initial stages of cosmic evolution (e.g., \\cite{Rees}; \\cite{Kronberg94}; and also see \\cite{Grasso} for a recent review). Recently, using the cosmological, magneto-hydrodynamic simulations of galaxy clusters, Dolag, Bartelmann, \\& Lesch (1999) have quantitatively shown that, if starting with the initial magnetic field of $\\sim10^{-9}$G strength, the final intracluster field can be amplified by the gravitationally induced collapsing motions to the observed strength of $\\sim\\mu$G irrespective of uniform or chaotic initial field configurations motivated by scenarios of the primordial or galactic-wind induced initial seed fields on Mpc scales, respectively. It is undoubtedly clear that the knowledge of the intracluster magnetic field will lead to a more complete understanding of the physical conditions of the intracluster medium and of the possible dynamical role of magnetic fields (e.g., \\cite{Loeb}; \\cite{Vikhlinin}). Furthermore, the magnetic field should play an essential role to the nonthermal processes such as the synchrotron and high energy radiations and possibly the cosmic ray production (\\cite{LoebNat}; \\cite{Totani}; \\cite{Waxman}). Hence, it is strongly desirable to perform further theoretical and observational investigations on these issues in more detail. In this paper we investigate how the Faraday rotation effect due to the intracluster magnetic field causes a secondary effect on the CMB polarization fields. While Thomson scattering of temperature anisotropies on the last scattering surface generates linearly polarized radiation at the decoupling epoch (\\cite{Kosowsky}; \\cite{Kami},b; \\cite{ZaldPol}; \\cite{HuWhite},b), the plane of linear polarization should be rotated to some extent when the CMB radiation propagates through the magnetized intracluster medium at a low redshift. The great advantage is that, since the CMB polarization field is a continuously varying field on the sky, we can measure an angular profile of the rotation measure in a cluster in principle, which cannot be achieved by any other means. Furthermore, thanks to the frequency dependence of the Faraday rotation, one will be able to separate this effect from the primary signal and other secondary signals such as that induced by the gravitational lensing effect due to the large-scale structure (\\cite{ZS98}). In this paper, based on the Press-Schechter theory for the cluster abundance under the cold dark matter (CDM) structure formation scenario and a plausible isothermal $\\beta$-model of the intracluster gas distribution, we compute angular power spectra of the CMB polarization fields including the Faraday rotation mixing effect caused by clusters at low redshifts. As for the unknown field configuration, we assume the uniform field of $\\sim 0.1\\mu{\\rm G}$ strength across a cluster that is consistent with the recent Faraday rotation measurements (Clarke et al. 2001). This model allows us to estimate the impact of this effect in the simplest way. Our study thus proposes a new empirical opportunity to measure or constrain the intracluster magnetic field in the average sense that properties of the magnetic field could be extracted through changes of the statistical quantities, CMB angular power spectra. So far previous works have focused mainly on investigations of effects of the primordial magnetic fields on the CMB temperature and polarization anisotropies (\\cite{KosoLoeb}; \\cite{Adams}; \\cite{Scan}; \\cite{Harari}; \\cite{Seshadri}; \\cite{Mack}; and see also \\cite{Grasso} for a review). Both the amplitudes of the cosmological primordial magnetic field and the mean baryonic density rapidly increase with redshift scaled as $(1+z)^2$ and $(1+z)^3$, respectively. The dominant contribution to the Faraday rotation effect is therefore imprinted before decoupling ($z\\simgt 10^3$), unless the structure formation at low redshifts causes a significant amplification of the magnetic field. Kosowsky, \\& Loeb (1996) found that, if the current value of the primordial field is of the order of $10^{-9}$G on Mpc scales corresponding to the microgauss galactic field under the adiabatic compression, the effect on the CMB polarization is potentially measurable by satellite missions {\\em MAP} and {\\em Planck Surveyor} (see also \\cite{Scan}; \\cite{Harari}). It is also shown that such a stochastic magnetic field in the early universe affects the temperature fluctuations through the induced metric vector perturbations and thus current measurements can put an upper limit of $\\simlt 10^{-11}$G on the current field amplitude (\\cite{Mack}). This paper is organized as follows. In Section \\ref{RM} we first construct a model to describe the magnetized hot plasma in a cluster by assuming the uniform magnetic field configuration for simplicity. We then derive the angular power spectrum of the Faraday rotation angle based on the Press-Schechter description of the cluster formation and the $\\beta$-model of the gas distribution. In Section \\ref{Formalism}, we present a formalism for calculating the angular power spectra of the CMB polarization fields including the Faraday rotation mixing effect. Section \\ref{results} presents the results for the currently favored CDM models. In Section \\ref{Disc} we briefly present the discussions and summary. Unless stated explicitly, we assume the favored $\\Lambda$CDM cosmological model with $\\Omega_{\\rm m0}=0.3$, $\\Omega_{\\lambda0}=0.7$, $\\Omega_{b0}=0.05$, $h=0.7$, and $\\sigma_8=1.0$ as supported from observations of CMB anisotropies and large-scale structure (e.g. \\cite{boom2}), where $\\Omega_{m0}$, $\\Omega_{b0}$, and $\\Omega_{\\lambda0}$ are the present-day density parameters of non-relativistic matter, baryon and the cosmological constant, respectively, $h$ is the Hubble parameter and $\\sigma_8$ denotes the rms mass fluctuations of a sphere of $8h^{-1}$Mpc radius. We also use the $c=1$ unit for the speed of light. ", + "conclusions": "\\label{Disc} In this paper, we have investigated the secondary effect on the CMB polarization fields induced by the Faraday rotation effect of the magnetic field in an intracluster hot plasma at a low redshift. To illustrate the impact of this effect in the simplest way, we employed the simple model that the magnetic field has a uniform field strength such as $0.2~\\mu{\\rm G}$ across a cluster universally when the cluster formed, which is consistent with observations of the rotation measures in clusters (Clarke et al. 2001). As shown in Figure \\ref{fig:averm}, our model can roughly reproduce the observed scatter of the rotation measures as a function of the distance from the cluster center (Kim et al. 1991; Clarke et al. 2001). In this sense, since the Faraday rotation effect on the angular power spectra of CMB polarizations comes from the second moments of the rotation measure angle (see equation (\\ref{eqn:rmclB})), it is likely that our model can at least estimate a correct magnitude of the effect, even though detailed shapes of those spectra could be different for a realistic intracluster magnetic field as discussed below. We showed that the parity-odd $B$-mode angular power spectrum is generated on the observed sky by the Faraday rotation mixing effect as a new qualitative feature, when if the primary CMB anisotropies includes the parity-even $E$-mode only such as predicted by the standard inflation-motivated scenarios with pure primordial scalar perturbations. We estimated that the generated $B$-mode power spectrum has a peak with amplitude of $\\sqrt{l^2C_{lB}/2\\pi}\\sim 0.1~\\mu{\\rm K} (B_0/0.1~\\mu{\\rm G}) (\\nu_0/10~{\\rm GHz})^{-2}$ at $l\\approx 1000$ under the plausible scenario of the cluster formation in the currently favored $\\Lambda$CDM model. It was also shown that the lowest frequency $30{\\rm GHz}$ channel of Planck can be used to set an upper limit of $B_0\\simlt 1.5\\mu{\\rm G}$ for a uniform component of the intracluster magnetic field. Detection or even null detection of the predicted $B$-type polarization will therefore be a new empirical tool to provide a calibration of an uncertain magnetic field in a cluster, combined with measurements of the gas distribution from the X-ray or SZ datasince. It is known that there are other secondary sources on the CMB polarization fields caused in the low redshift universe, and these nonlinear effects generally induce the $B$-type polarization pattern from the coupling with the primary $E$-mode. The main source is the gravitational lensing effect of the large-scale structure, leading to the $B$-type polarization that has a peak of $\\sim 0.3~\\mu{\\rm K}$ at $l\\approx 1000$ and then a power of $\\simlt 0.01~\\mu{\\rm K}$ at $l\\simgt 5000$ for the simliar $\\Lambda$CDM model as considered in this paper (\\cite{ZS98}). One of authors (N.S.) has quantitatvely shown that the secondary effect due to the peculiar motion of ionized medium in the large-scale structure also generates the $B$-mode polarization with $\\sim 0.01~\\mu{K}$ at $100\\simlt l \\simlt 10^4$ for a realistic patchy reionization model of the universe (\\cite{Liu2001}). From these results, if the uniform magnetic field has a larger strength of $B\\sim 0.5~\\mu{\\rm G}$ as a common feature of clusters and we observe at lower frequencies such as $\\nu_0\\simlt 5~{\\rm GHz}$, the generated $B$-mode amplitude due to the Faraday rotation effect would be comparable with the primary polarization amplitude and larger than the other secondary signals at $l\\simgt1000$, although observations at such low frequencies would suffer from large foreground contamination of the syncrhrotron radiation from our galaxy (\\cite{Tegmark}). Even if the uniform magnetic field is not as strong as $0.5~\\mu{\\rm G}$, the sensitive multi low-frequency measurements will allow one to separate this signal from primary and those other secondary signals in principle thanks to the frequency dependence of $\\propto \\nu_0^{-4}$ on $C_{Bl}$, which is analogous to measurements of secondary temperature fluctuations by the SZ effect. There are possible contributions to affect our results that we have ignored in this paper. The intracluster tangled magnetic fields, as implied by the rotation measures of some sources inside the Coma cluster (\\cite{Kim90}), would provide contributions to our results in addition to the contribution due to the uniform field. However, the coherent scale of tangled fields in a cluster is still unknown. The recent complete measurements of the rotation measures (Clarke et al. 2001) imply the $1-10~\\mu{\\rm G}$ strength of the tangled field under the assumption of the coherent scale of $10-100~{\\rm kpc}$. It is also known that the strong radio halos embedded within a (cooling flow) cluster are pervaded by stronger magnetic fields on smaller scales such as $\\sim 50~{\\mu G}$ for Hydra A (\\cite{Taylor}; \\cite{Taylor01}). Although the simplest tangled-cell model with a constant coherence length $l$ has often been considered in the literature, the possible important mechanism to affect our results is the deporlarization effect (\\cite{Tribble}), leading to the cancellations to some extent between the Faraday rotations in independent cells with random orientations of the magnetic fields within a cluster. The finite beam size will also lead to the artificial cancellations between the Faraday rotations of cells covered within one beam. From these considerations, the depolarization effect may lead to an intuitive result that a feature of the $B$-mode spectrum can be sensitive to the characteristic angular scale, $l_{c}$, originating from the coherent length. This investigation is now in progress, and will be presented elsewhere (Ohno et al. 2001). Recently, it has been suggested that outflow or jet acitivities of quasars or black holes leave behind an expanding magnetized bubble in the itergalactic medium (IGM) resulting in the possibility that the IGM is filled by magnetic fields to some extent (\\cite{Furlan}; \\cite{Kronberg01}). Even thgouh the IGM field has a smaller strength than $\\sim 10^{-9}{\\rm ~G}$ on ${\\rm Mpc}$ scales, a large volume coverage of the ionized plasma might lead to non-negligible contribution connected to the reionization history of the universe. Finally, we comment on implications of our results to the nonthermal processes in a cluster. The observations of diffuse synchrotron halos and nonthermal hard X-ray emission have implied the existence of nonthermal relativistic electrons in a cluster (e.g., \\cite{Rephaeli99}). These electrons are produced by the shock acceleration in an intracluster medium, and it will allow us to derive additional information on the physical conditions of the intracluster medium environment, which cannot be obtained from the thermal plasma emission only. The intracluster magnetic field should then play an important role to these nonthermal processes. In particular, it is being recognized that the nonthermal process can be a unique probe of measuring the dynamically forming clusters before thermalization of the intracluster gas. Recently, Waxman \\& Loeb (2000) pointed out that the synchrotron radiation from those forming clusters could be foreground sources of the CMB temperature fluctuations at low frequency observations such as $10{\\rm ~GHz}$. Those clusters could also produce the high energy gamma ray emission through the inverse Compton scattering of CMB photons by relativistic electrons and thus be identified as {\\it gamma ray clusters} by the future sensitive telescope {\\em GLAST}, which are difficult to detect by X-ray or optical surveys because of their extended angular size of about $\\sim 1^\\circ$ (\\cite{Totani}). Based on these ideas, it will be very interesting to investigate cross-correlations between the synchrotron radiation or the gamma ray sources and the Faraday rotation signal in CMB polarization map. This investigation is expected to provide complementary constraints on the nonthermal physical processes in the forming clusters." + }, + "0112/astro-ph0112138_arXiv.txt": { + "abstract": "{We introduce analytical expressions for a pseudo fully analytical elliptical projected Navarro, Frenk \\& White (NFW) mass profile to be used in lensing equations. We propose a formalism that incorporates the ellipticity into the expression for the lens potential, producing a pseudo-elliptical mass distribution. This approach can be implemented to any circular mass profile for which the projected mass profile $\\Sigma(r)$ and the deflection angle profile $\\alpha(r)$ both have analytical expressions; however the potential does not necessarily need to take an analytical form. We apply this new formalism to the NFW mass distribution and study how well this pseudo-elliptical NFW model describes an elliptical mass distribution. We conclude that the pseudo-elliptical NFW model is a good description of elliptical mass distributions provided that the ellipticity of the projected mass distribution is $\\lesssim 0.4$, although with a slightly boxy distribution. ", + "introduction": "Cosmological N-body simulations of cluster formation \\citep{NFW} indicate the existence of a universal density profile for dark matter halos, independent of their mass, power spectrum of initial fluctuations or cosmological parameters. For this so-called NFW profile, the density increases near the centre with a shallower slope than an isothermal profile, while it steepens gradually outward and becomes steeper than isothermal far from the centre. Its analytic expression is given by \\begin{equation} \\rho(r)=\\frac{\\rho_{\\mathrm c}}{(r/r_s) (1+r/r_s)^2} \\label{rho_nfw} \\end{equation} where $\\rho_{\\mathrm c}$ is a characteristic density and $r_s$ a scale radius. Recent higher-resolution simulations \\citep[e.g.][]{Moore,Ghigna} advocate a steeper central cusp of $\\rho\\propto r^{-1.4}$. Attempts to constrain the inner slope of the density profile with high resolution observations of luminosity profiles \\citep{Faber} seems to confirm a central cusp ($\\rho\\propto r^{-1}$), rather than a core radius for massive galaxies. On larger scales, \\citet{Smith} used gravitational lensing to constrain the density profile of A\\,383, a massive galaxy cluster at $z=0.19$, finding a logarithmic slope of $\\sim-1.3$. Robust interpretation of these observational results is complicated by several factors, including the absence of baryons from high resolution numerical simulations, systematic uncertainties in the lens models arising from parametrisation of the mass distribution, and the need to use elliptical mass distributions to fit observed multiple image systems. Gravitational lensing is an ideal tool to constrain the radial structure of collapsed halos such as galaxies and clusters of galaxies \\citep{Smith}. However, lensing is only sensitive to the projected mass distribution, and elliptical mass distributions are needed to match the multiple images observed in both galaxy and cluster lens systems \\citep{Kneib01}. In response to the debate regarding the inner slope of the density profile, \\citet{Munoz} introduced a general set of ellipsoidal lens models with $\\rho\\propto r^{-\\gamma}$ as $r \\rightarrow 0$ and $\\rho\\propto r^{-n}$ at large radius. However, as there are no general analytic expressions for cusped ellipsoidal mass models, the deflections and magnifications are calculated numerically. They applied their model to the gravitational lens APM 08279+5255 and found a very shallow cusp ($\\gamma\\lesssim0.4$). In contrast, for B~1933+503, they found that a steep density cusp ($1.6\\lesssim\\gamma\\lesssim2.0$) is favoured. To avoid expensive numerical integration, \\citet{Barkana} suggested an alternative solution. For a softened power-law elliptical mass distribution, it is possible to approximate the integrand so that the integration can be done analytically. Therefore, for this flat core model, the deflection can be then calculated to high accuracy. In this paper we propose a new way to introduce ellipticity in lensing model in a fully analytical way, and we discuss in detail the recipe and limit of the model for the NFW mass distribution. In Sect.~2, we briefly discuss spherical NFW lens models. Then we present, in Sect.~3, a general pseudo-elliptical formalism that incorporates the ellipticity in the expression of the lens potential if this is known, or anyway of the deflection angle. In Sect.~4, we apply this formalism to the NFW profile and study the departure of this model from an elliptical NFW mass model. Finally, in Sect.~5 we discuss prospects for the application of this new formalism. ", + "conclusions": "We propose a simple new formalism that introduces the ellipticity into the lens potential/deflection-angle of a circular mass model. The method can be applied when the lens potential or/and the deviation angle takes an analytical form. Then for radial mass profiles for which the 2D surface density $\\Sigma$ also has an analytical expression, this formalism gives analytical expressions of a pseudo-elliptical mass distribution for the deviation angle, the projected mass density, the convergence and shear. Whatever the form of the mass distribution, the elliptical parameter $\\epsilon$ is simply expressed as a function of the ellipticity of the potential. This is particularly helpful in getting some insight on the physical meaning of this parameter. We have applied this formalism to the NFW profile and estimated the range of ellipticity ($\\epsilon\\lesssim0.25$, or $\\epsilon_\\Sigma\\lesssim0.4$) for which this model is a good description of elliptical mass distributions and thus can be reliably applied to observational data. To derive these limits, we introduced a particular fit for elliptical-like profiles, that can be useful in similar cases. Our proposed method is particularly useful when it is essential to quickly calculate the potential, the deflection angle and magnification of many images and/or many mass clumps. This is particularly important when using {\\it inverse} methods (such as maximum likelihood) to investigate galaxy-galaxy lensing in the field or in clusters of galaxies, or to compute time delays." + }, + "0112/astro-ph0112335.txt": { + "abstract": "We present 0.9---2.5$\\mu$m spectroscopy with R$\\sim$800 and 1.12-1.22$\\mu$m spectroscopy with R$\\sim$5800 for the M dwarfs Gl 229A and LHS 102A, and for the L dwarf LHS 102B. We also report $IZJHKL^{\\prime}$ photometry for both components of the LHS~102 system, and $L^{\\prime}$ photometry for Gl~229A. The data are combined with previously published spectroscopy and photometry to produce flux distributions for each component of the kinematically old disk M/L--dwarf binary system LHS~102 and the kinematically young disk M/T--dwarf binary system Gliese~229. The data are analyzed using synthetic spectra generated by the latest ``AMES-dusty'' and ``AMES-cond'' models by Allard \\& Hauschildt. Although the models are not able to reproduce the overall slope of the infrared flux distribution of the L dwarf, most likely due to the treatment of dust in the photosphere, the data for the M dwarfs and the T dwarf are well matched. We find that the Gl~229 system is metal--poor despite having kinematics of the young disk, and that the LHS~102 system has solar metallicity. The observed luminosities and derived temperatures and gravities are consistent with evolutionary model predictions if the Gl~229 system is very young (age $\\sim 30\\,$Myr) with masses (A,B) of (0.38,$\\ga$ 0.007)M$_{\\sun}$, and the LHS~102 system is older, aged 1---10~Gyr with masses (A,B) of (0.19,0.07)M$_{\\sun}$. ", + "introduction": "In the last few years there have been dramatic developments in the study of very low mass stars and brown dwarfs --- where the latter are defined to be objects with mass below that required for stable hydrogen burning. Sky surveys have found a significant population of low mass dwarfs cooler than M--dwarfs. The first population to be identified was the L--dwarf class, distinguished from the M--dwarfs by weakening VO and TiO absorption features in the red, and by stronger alkaline and H$_2$O absorption features in the red and near--infrared \\cite[e.g.][]{d97,k99,mar99}. The L--dwarfs cover an effective temperature range of about 2200~K to 1400~K \\cite[e.g.][]{l01} and the prototype of this class is GD~165B, discovered by \\cite{bz88} as a red companion to a hot white dwarf. Several objects even cooler than the L--dwarfs were discovered in 1999 by the Sloan Digital Sky Survey and the 2 Micron All--Sky Survey \\cite[]{str99,bur99}. They are very similar to Gliese 229B, the extremely low--mass companion to an M--dwarf found by \\cite{nak95}, and are distinguished by the presence of CH$_4$ absorption in the near--infrared H and K bands. These objects, known as T--dwarfs, have $\\Teff \\sim$1300---800~K and tens of T--dwarfs are now known \\cite[see][and references therein]{g01,b01}. In this paper we present an analysis of the low--mass binary systems LHS~102 and Gliese (Gl)~229. The primaries of these two systems are M--dwarfs which are in the proper motion catalogues by \\cite{luy79} and \\cite{gj91}. The secondaries were only recently detected. Gl~229B, the prototypical T--dwarf, was discussed above; LHS~102B is an L dwarf discovered by \\cite{go99}. We present new UKIRT infrared photometry and spectroscopy of these objects in \\S 2, which are analyzed using the latest models by Allard \\& Hauschildt \\citep*{ng-hot}, described in \\S 3. The results of the model comparison are given in \\S 4, the implications for age and mass of each component are discussed in \\S 5 and our conclusions given in \\S 6. These two systems span the newly defined low--mass spectral classes and, being binaries at known distances where each component has presumably the same chemical composition and age, are potentially very useful for constraining atmospheric and evolutionary models of ultracool dwarfs. ", + "conclusions": "In this paper we have presented the results of the analysis of two binary systems, one (LHS 102) consisting of an M/L dwarf pair and one (Gl 229) being an M/T dwarf system. The M dwarf primaries show that the best mixing length for effective temperatures above $3000\\K$ is somewhere around $1.5$ to $2.0$, which is consistent with recent hydrodynamical models \\cite[]{ludwig-pap}. The LHS 102 system is best fit with solar abundance models. LHS 102B is an L dwarf with an effective temperature around $1900\\K$ that is best fit with ``AMES-dusty'' models that include the effects of dust formation and opacity on the atmospheric structure and the emitted spectrum. More work is required on the detailed treatment of photospheric dust to produce a good match to the entire observed flux distribution; such work is in progress. LHS 102B is found to be just below the stellar mass limit. The results for Gl 229 indicate that this system is metal deficient with $\\mh \\approx -0.5$. We have determined an age for the Gl 229 system of about 30 Myr, constrained primarily by the observed luminosity and derived effective temperature of the A component. This young age is consistent with the low surface gravity derived for Gl 229B \\cite[see for example Figure 9 of][]{bur97}, and translates to a very low mass for this T dwarf of $\\ga$~7 Jupiter masses. A better mass determination requires metal-poor evolutionary models of brown dwarfs, which are not currently available." + }, + "0112/astro-ph0112295_arXiv.txt": { + "abstract": "I describe some of the current challenges in galaxy formation theory with applications to formation of disks and of spheroids. Forthcoming deep surveys of galaxies with Keck and VLT will provide high quality spectra of $\\sim 10^5$ galaxies that will probe stellar populations and star formation rates at redshift unity. This will help refine our phenomenological knowledge of galaxy evolution and enable robust predictions to be developed for future breakthroughs in understanding galaxy formation at high redshift that are anticipated with NGST and with the proposed new generation of 30 metre-class telescopes. ", + "introduction": "Galaxy formation is a complex process, involving both gravity and hydrodynamics, and can be complicated by such ingredients as turbulence and astrochemistry. The disks and the spheroids of galaxies have undergone distinct. although not necessarily uncoupled, histories. The physics of disk formation has made considerable progress, in no small part due to the pioneering review by Ken Freeman in {\\it Stars and Stellar Systems, Volume IX} that assembled diverse observational and theoretical aspects together for the first time. Spheroid formation is in a less satisfactory state, in part because spheroids are old and so their formation occurred long ago, and there are correspondingly few direct clues. We do not yet have an adequate understanding of either disk or spheroid formation. There is a simple reason for this predicament. We have no fundamental theory of star formation: the best we can do even in nearby regions of star fomation is to assemble phenomenologically-motivated arguments and laws. When phenomenology is sparse as in the early universe, all bets are off as to the scalability of current epoch theory to the past. Of course, the lack of a robust theory has never deterred theorists, and in this talk I will highlight some of the key issues currently confronting cosmologists. ", + "conclusions": "" + }, + "0112/astro-ph0112540_arXiv.txt": { + "abstract": " ", + "introduction": "Clusters of galaxies are of wide interest in the astronomical community both as laboratories where baryons and dark matter can be conveniently studied and, through their statistical properties at various look-backtimes, as tracers of cosmic structure and evolution. The Local Organising Committee wisely narrowed the topics of this meeting to four broad themes and so I will organise my summary remarks under these headings. First I think I should address the plea \"Why bother to study clusters?\", raised provocatively by Colin Norman at one stage during the proceedings\\footnote{Fortunately, I have forgotten to which speaker he made this remark!}. We have learnt enough about clusters and the growth of structure to know that two of the traditional motivations require much more careful consideration. For years on time allocation committees, I read observers repeat a {\\em mantra} in the first sentence of their proposals: {\\em ``Clusters of galaxies represent the largest bound structures in the Universe..\"} in justification of their role as tracers of large scale structure. A number of speakers at this meeting reminded us that clusters no longer uniquely occupy this role. We have other tracers of large scale structure and, moreover, it seems we need to be much clearer of what exactly we mean by a cluster before we can convincingly use them as cosmological probes. A second, traditional, motive for studying clusters has been observational convenience, e.g. in studying constituent populations such as galaxies at various look-back times. At first sight, it's an attractive proposition for an observer to study a few rich clusters at various redshifts each containing hundreds of accessibly-luminous galaxies and then to ``join the results\" in a timeline to make some evolutionary claim. But, at most redshifts of interest, we can expect to find a wide range of overdensities represented (groups, merging systems, virialised clusters). Put simply, rich cluster A at $z$=1.2 is unlikely to evolve into rich cluster B at $z$=0.5 and neither may necessarily become a present-day Coma. Both these worries indicate a high degree of rigour is needed in using clusters. We need very large samples spanning ranges of mass at any redshift of interest and perhaps it would be foolish to adopt one single selection criterion for their study. A similar ``panchromatic\" theme has emerged over the past decade in understanding how to address the question of evolution using samples of galaxies.\\footnote{Witness the controversy surrounding the use of UV/optical and far-infrared probes of the cosmic star formation history. No single technique wins: several are required.} The motivation for studying clusters that emerged at this meeting focused broadly along the following themes: \\begin{itemize} \\item{} Testing gravitational instability by measuring the number density of massive clusters at high redshift, viz. $\\Phi(M_{cluster},z)$. \\item{} Breaking degeneracies in estimates of the the cosmological parameters by examining the local population of clusters as a probe of the mean mass density and the normalisation of the mass power spectrum, $\\Omega_M$ and $\\sigma_8$. \\item{} Verifying hierarchical structure and the nature of the dark matter and keeping those theorists in check who predict a universal mass profile with central cusps, $\\rho_M(r) \\propto r^{-1.5}$. \\item{} Determining the origin of the heating of the intracluster medium, its enrichment history and examining whether non-gravitational processes are involved. \\item{} Examining the history of spheroidal galaxies and role of the environment, for example in understanding the origin of the morphology-density relation and the destiny of the infalling component of gas-rich field galaxies. \\end{itemize} Let me make a disclaimer in what follows. I am not an expert in any of the areas I was asked to summarise and so I submit these concluding remarks only as someone who tried to listen carefully to most talks, dutifully avoiding the lunchtime mountain hikes to make sense of what I heard\\footnote{I apologise if I missed some talks but I tried to secure powerpoint files or transparencies for most of those.}. With $\\simeq$70 talks and $\\simeq$80 posters, inevitably I have had to be very selective in discussing results. If one could summarise the meeting in two paragraphs, I would say the following: \\begin{itemize} \\item{} There is an explosion of terrific data, from large surveys (X-ray and ground-based) which offer qualitatively new ways in which we find and do statistical studies of clusters, and also in {\\em resolved} data within clusters which opens up new opportunities for understanding the detailed astrophysics of dense environments. \\item{} The subject is moving from exploratory surveying to detailed astrophysics. As part of this ``growing up\" there is a need to admit defeat on some of the old methods and to embrace new ones, particularly on the statistical questions. For many years cluster workers had something of a monopoly in the study of early galaxies and large scale structure, but it is time to take advantage of other datasets being delivered and view cluster astrophysics as only one part of a larger body of information. \\end{itemize} ", + "conclusions": "" + }, + "0112/astro-ph0112226_arXiv.txt": { + "abstract": "This article gives a brief qualitative description of the possible evolution of the early Universe between the end of an inflationary epoch and the end of Big Bang nucleosynthesis. After a general introduction, establishing the minimum requirements cosmologists impose on this cosmic evolutionary phase, namely, successful baryogenesis, the production of cosmic dark matter, and successful light-element nucleosynthesis, a more detailed discussion on some recent developments follows. This latter includes the physics of preheating, the putative production of (alternative) dark matter, and the current status of Big Bang nucleosynthesis. \\footnote{Article based on a talk presented at ``The Early Universe and Cosmological Observations: a Critical Review'', Cape Town, July 2001} ", + "introduction": "I was asked by the organizers of the conference to present a review of processes occurring in the early Universe, between the end of a cosmological inflationary period and the end of Big Bang nucleosynthesis. Due to the multitude of possible processes and phenomena which may have occurred over such a long evolutionary interval, and the plethora of scientific work which exists on the subject, accomplishing such a task on a few pages is a difficult undertaking and bound to be done superficially. Nevertheless, I realize that such a mini-review of reviews (i.e. frequent citing of other review articles) may be of some aide to scientists outside of the community and provide an initial starting point for further study. It is in this way I like the article to be understood. My apology goes to all scientists whose work has not been mentioned explicitly. The article is split into two parts. The first part (sections 2 and 3) contains a very general account of the evolution of the early Universe and the processes which must have occurred to create an ``acceptable'' Universe. In the second part (sections 4 - 9) more emphasis is laid on some recent developments and advances in the field. Note that the article does not consider cosmological phenomena due to the (putative) existence of extra spatial dimensions (cf. David Wands, these proceedings), as the study of such is still in it's infancy. \\begin{figure}[t] \\centering \\leavevmode\\epsfysize=12cm \\epsfbox{slide1.ps}\\\\ \\end{figure} ", + "conclusions": "In conclusion, let me just make a short remark. Early Universe cosmology in the eighties discovered the virtues of a (Hubble) damping term in the cosmological equation of motion for bosonic fields. This discovery led to the invention of inflation. In the nineties the virtues (and problems) of terms leading to instabilities in the equation of motion for bosonic fields in the Universe were discovered. This revelation led to the possibility of cosmological preheating and formation of Q-balls. Inspection of the remaining terms in the equations of motion for a bosonic field in the Universe may give, the educated reader, an outlook on possible developments in the first decade of this millennium. \\ack The author thanks the organizers of the conference for invitation and financial support, and the opportunity to visit South Africa. I am also grateful to my girlfriend Nanci, who ``forced'' me to be a responsible scientist and write this article." + }, + "0112/astro-ph0112360_arXiv.txt": { + "abstract": "We report the serendipitous detection of two previously unreported pulsars from the direction of the Small Magellanic Cloud, with periods of 16.6 and 25.5 seconds. The detections are based on archival PCA data from the {\\it Rossi X-Ray Timing Explorer} (RXTE). The observation leading to these detections occurred in September 2000 extending over 2.1 days with an exposure of 121ks. A possible identification of the 16.6 pulsar with an X-ray source (RX J 0051.8-7310) seen by both ROSAT and ASCA imaging X-ray satellites is presented. ", + "introduction": "In the course of a survey of archival data from the Proportional Counter Array (PCA) of the {\\it Rossi X-Ray Timing Explorer}(RXTE: Bradt, Rothschild, \\& Swank 1993), we have found evidence for two previously unreported pulsars with periods of 16.6 and 25.5 seconds.. The evidence is based on observations taken 2000 September 13-15, with the PCA field of view, $\\sim$1.0$^\\circ$ FWHM, (Jahoda et al. 1996) viewing an area near the southwestern edge of the Small Magellanic Cloud, centered on 00$^h$ 50$^m$ 44.64$^s$, -73$^\\circ$ 16$^{'}$ 04.8$^{``}$. The observation in question extended for 2.1 days, with minimal interruptions, leading to 68\\% of the time devoted to source region coverage. Such extended ``dense'' observations are well-suited for sensitive searches for periodic phenomena from relatively faint sources. The PCA observations used the so-called Good Xenon data mode, in which the arrival of each photon at the detector is time tagged to an accuracy of better than 1$\\mu$s. For the timing analysis the electrical pulses were required to originate in either one of the top two layers (of three) of the PCA, with pulse heights in the X-ray energy range 2 to 15 keV. Such an energy and layer selection is effective in improving signal-to-noise for pulsar detection. The times of arrival were corrected to the barycenter of the solar system and a discrete Fourier transform (fft) of these times was performed. The portion of this transform from 0.02 to 0.2 Hz is shown in Figure 1. \\begin{figure} \\figurenum{1} \\plotone{f1.eps} \\caption{FFT of the RXTE Observation} \\end{figure} There are four distinct frequencies and one second harmonic which are evident in the spectrum of Figure 1. The frequencies, corresponding periods, and Fourier power (normalized to 1) are listed in Table 1. Two of the periods are consistent with those that have been previously reported from this region of the SMC. The 8.9 s period is identified with the same period pulsation from the Be transient RX J0051.8-7231 (Israel et al. 1997). The 15.7 period may be identified with the reported 15.3 s period from the SMC transient RX J0052.1-7319 (Finger et al. 2001). The observations reported by Finger et al. for this object occurred 3.8 years earlier than the RXTE observation of Figure 1. The difference in period, 2.5\\%, can be accounted for by a spin-down episode(s) consistent with spin histories of other Be transients (Bildsten et al. 1997), therefore we regard this identification as reasonable. For the remaining periods, 16.6 and 25.5 s, we have found no previous reports covering this portion of the sky. \\begin{deluxetable}{rrr} \\tablecolumns{3} \\tablecaption{List of Frequencies with Power $>$20} \\tablewidth{0pt} \\tablehead{\\colhead{Frequency(Hz)} & \\colhead{Nomalized Power} & \\colhead{Period (s)}} \\startdata 0.039230 & 48. & 25.5 \\\\ 0.060339 & 196. & 16.6 \\\\ 0.063861 & 54. & 15.7 \\\\ 0.112401 & 58. & 8.90 \\\\ 0.120685 & 22. & 16.6 (2Harm) \\enddata \\end{deluxetable} In what follows we present the lightcurves for both the 16.6 and the 25.5 s pulsars and discuss evidence for a possible identification of the 16.6 s pulsar with a ROSAT source, RX J0051.8-7310. In the concluding section we discuss the luminosities of sources. If these pulsars are associated with the SMC, as is likely, their discovery further accentuates the remarkable overdensity of binary pulsars in the SMC relative to our own galaxy. ", + "conclusions": "For the 25.5s pulsar we derive a pulsed flux value of 6.9$\\pm$0.8$\\times$10$^{-12}$ (2-15 keV) under model assumptions of a power-law spectrum with an energy index of -2.0 and a hydrogen column density of 1.0$\\times$10$^{21}$atoms/cm$^2$. Under the same assumptions the 16.6 pulsar flux is 8.5$\\pm$1.1$\\times$10$^{-12}$ (2-15 keV). At the distance of the SMC (nominally 60 kpc) these pulsed fluxes correspond to isotropic luminosities of 3.0 and 3.7$\\times$10$^{36}$ ergs/s. These fluxes are a few percent of the Eddington luminosity, 2$\\times$10$^{38}$ ergs,s, for a 1.4 solar mass object. Two arguments can be made to support the idea that at least one if not both of these pulsars are binary. Isolated pulsars with periods greater than 10s have yet to be seen. Also there is a significant non-zero period derivative for the 16.6 s pulsar with a time-scale of $\\sim$ 50 years. Further analysis of existing and future RXTE observations may be able to constrain possible orbital parameters for this object. If these pulsars are binaries in the SMC, their discovery further accentuates the dramatic difference between the SMC and our Galaxy with regard to the population of such systems. This fact has been noted by several authors (Schmidtke et al. 1999, Yokogawa et al. 2000). In a recent compilation of known X-ray pulsars $^4$ there are 18 X-ray pulsar binaries listed for the SMC all of which are either high-mass or transient and therefore likely to be high mass systems. We cannot say if these newly discovered pulsars are high or low-mass. However, if they are high mass, then this increases the number of such systems in the SMC to 20. For the Galaxy the corresponding number is 40. Therefore, using a mass ratio of the SMC to the Galaxy of 1/100, this suggests that such systems are over-abundant by a factor of $\\sim$50 relative to the Galaxy. This simple analysis ignores important issues regarding the uncertain coverage of the Galaxy for transient X-ray binaries versus the rerlatively complete coverage of the SMC. It also ignores differences in X-ray absorption effects. Nevertheless, pending careful analysis of such issues, there appears to be a significant overabundance of high mass binaries in the SMC relative to the Galaxy. Since high mass X-ray binaries have lifetimes which are $\\sim10^{-3}$ the age of the Galaxy and possibly the SMC, this difference between the SMC and the Galaxy may point to a rather recent outburst of star-formation in the SMC within the past $\\sim10^{7}$ years. Yokogawa et al. (2000) reach a similar conclusion. $^4$ http://gammaray.msfc.nasa.gov/batse/pulsar/asm.pulsars.html" + }, + "0112/astro-ph0112156_arXiv.txt": { + "abstract": "The physics of gamma-ray bursts (GRBs) and their offsets from the centers of their host galaxies are used to investigate the evolutionary state of their progenitors, motivated by the popular idea that GRBs are linked with the cataclysmic collapse of massive stars. We suggest that GRB progenitors in the inner and outer regions of hosts may be intrinsically different: outer bursts appear to have systematically greater isotropic equivalent energies (or narrower jets). This may provide an interesting clue to the nature of GRBs, and could reflect a relation between metallicity and the evolution of GRB progenitors. If true, then this offset--isotropic luminosity correlation is a strong argument for a collapsar origin of long-duration GRBs. ", + "introduction": "One can understand the dynamics of GRB afterglows simply, independent of uncertainties about their progenitors, using the relativistic generalization of the theory of supernova remnants. The basic model for GRB hydrodynamics is of a relativistic blast wave that expands into the surrounding interstellar medium (ISM; M\\'esz\\'aros \\& Rees 1997), decelerates on contact with the ambient matter, and leads to a predictable radiative spectrum with a characteristic power-law decline. The study of GRB afterglows has provided confirmation of relativistic source expansion (Piran 1999; M\\'esz\\'aros 2001). The energy source of the fireball is assumed to be a cataclysmic event, either a compact stellar merger (Lattimer \\& Schramm 1976; Eichler et~al.\\ 1989) or the collapse of a massive star (Woosley 1993; Paczynski 1998; MacFadyen \\& Woosley 1999, hereafter MW99). Evidence is accumulating that GRBs are intimately linked with the deaths of massive stars. For the long-burst afterglows localized so far, the host galaxies show signs of the ongoing star formation activity necessary for the presence of young, massive progenitor stars (Kulkarni$\\;$ et al. 1998; Fruchter et al. 1999; Berger et al. 2001). The physical properties of the afterglows, their locations in host galaxies (Bloom, Kulkarni \\& Djorgovski 2001b), iron line features (Piro et al.~2000; Amati et al.~2000), and evidence for supernova components several weeks after three bursts (GRB980326, Bloom et al. 1999a; GRB970228, Reichart 1999; GRB 000911, Lazzati et al. 2001) strongly support the idea that the most common GRBs are linked to the collapse of massive stars.\\\\ The circumburst medium provides a natural laboratory for studying GRBs. Stars that readily shed their envelopes have short jet-crossing times and are more likely to produce a GRB. Stars with less radiative mass loss retain a hydrogen envelope, in which a poorly collimated jet is likely to lose energy and fail to breaking out of the star (MW99). Finding useful diagnostics for the progenitors is simplified if the metallicity of and physical conditions in the local ISM influences the evolution of the progenitor. GRBs occur close to the birth sites of their short-lived progenitors, and so their evolution is likely to be affected only by local properties of the host galaxy. Here, we show that bursts located closer to the center of their parent galaxies have smaller isotropic equivalent energies (or broader jets), and so progenitors in inner and outer galactic locations may be intrinsically different. We suggest that this could be the outcome of abundance gradients in the host galaxy. We assume $H_0 = 65\\,\\, {\\rm km} \\, {\\rm s}^{-1} \\, {\\rm Mpc}^{-1}$, $\\Omega_{\\rm matter}=0.3$, and $\\Omega_{\\Lambda}=0.7$. ", + "conclusions": "We report a correlation between the isotropic equivalent energy of GRBs and their position offset from their host galaxies. This is possibly due to a dependence of the end point of massive stellar evolution on metallicity. If confirmed in further host observations, this correlation will both complicate interpretation of GRBs as tracers of cosmic star formation, and potentially allow a new probe of the astrophysics in high-redshift galaxies." + }, + "0112/astro-ph0112427_arXiv.txt": { + "abstract": "Dark-matter halos are the scaffolding around which galaxies and clusters are built. They form when the gravitational instability of primordial density fluctuations causes regions which are denser than average to slow their cosmic expansion, recollapse, and virialize. Objects as different in size and mass as dwarf spheroidal galaxies and galaxy clusters are predicted by the CDM model to have halos with a universal, self-similar equilibrium structure whose parameters are determined by the halo's total mass and collapse redshift. These latter two % are statistically correlated, however, since halos of the same mass form on average at the same epoch, with small-mass objects forming first and then merging hierarchically. The structural properties of dark-matter dominated halos of different masses, therefore, should reflect this statistical correlation, an imprint of the statistical properties of the primordial density fluctuations which formed them. Current data reveal these correlations, providing a fundamental test of the CDM model which probes the shape of the power spectrum of primordial density fluctuations and the cosmological background parameters. ", + "introduction": "We have developed an analytical model for the postcollapse equilibrium structure of virialized objects which condense out of a cosmological background universe, either matter-dominated or flat with a cosmological constant \\cite{SIR,ISb}. The model is based upon the assumption that cosmological halos form from the collapse and virialization of ``top-hat'' density perturbations and are spherical, isotropic, and isothermal. This leads to a unique, nonsingular TIS, a particular solution of the Lane-Emden equation (suitably modified when $\\Lambda\\neq0$). The size $r_t$ and velocity dispersion $\\sigma_V$ are unique functions of the mass and redshift of formation of the object for a given background universe. Our TIS density profile flattens to a constant central value, $\\rho_0$, which is roughly proportional to the critical density of the universe at the epoch of collapse, with a small core radius $r_0\\approx r_t/30$ (where $\\sigma_V^2=4\\pi G\\rho_0r_0^2$ and $r_0\\equiv r_{\\rm King}/3$, for the ``King radius'' $r_{\\rm King}$, defined by \\cite{BT}, p. 228). The density profiles for gas and dark matter are assumed to be the same (no bias), with gas temperature $T=\\mu m_p\\sigma_V^2/k_B$. \\begin{figure} \\centering \\begin{minipage}[c]{0.45\\textwidth} \\hspace{-2.2cm} \\includegraphics[width=2.5in]{iliev_F1.eps} \\label{profile} \\end{minipage} \\begin{minipage}[c]{0.4\\textwidth} \\centering \\begin{tabular}{@{}lcccc} Table 1 &&&\\\\ &SUS&SIS&TIS\\footnote{The two values refer to flat universe with $\\rm \\Lambda=0$ (left value) and $\\rm \\Omega_0=0.3$, $\\lambda_0=0.7$ (right value).}\\\\ \\hline $\\eta/\\eta_{\\rm SUS}$&1&0.833&1.11;1.07\\\\[1mm] $% {T}/{T_{\\rm SUS}}$&1&3&2.16;2.19\\\\[1mm] $\\displaystyle{{\\rho_0}/{\\rho_t}}$& 1&$\\infty$&514;530\\\\[1mm] $\\displaystyle{{\\langle\\rho\\rangle}/{\\rho_t}}$&1&3&3.73;3.68\\\\[1mm] $\\displaystyle{{r_t}/{r_0}}$& -- NA --&$\\infty$&29.4;30.04\\\\[1mm] ${\\Delta_c}/{\\Delta_{\\rm c,SUS}}$&1& 1.728&0.735;0.774\\\\[1mm] $K/|W|$&0.5&0.75&0.683;0.690 \\\\\\hline \\end{tabular} \\end{minipage} \\caption{(top) Density profile of TIS % in a matter-dominated universe. Radius $r$ is in units of $r_m$ - the top-hat radius at maximum expansion. Density $\\rho$ is in terms of the density $\\rho_{SUS}$ of the SUS approximation for the virialized, post-collapse top-hat. (bottom) Logarithmic slope of density profile.} \\end{figure} These TIS results differ from those of the more familiar approximations in which the virialized sphere resulting from a top-hat perturbation is assumed to be either the standard uniform sphere (SUS) or else a singular isothermal sphere (SIS). We summarize their comparison in Fig. 1 and Table 1, where $\\eta$ is the final radius of the virialized sphere in units of the top-hat radius $r_m$ at maximum expansion (i.e. $\\eta_{\\rm SUS}=0.5$), $\\rho_t\\equiv\\rho(r_t)$, $\\langle\\rho\\rangle$ is the average density of the virialized spheres, $\\Delta_c=\\langle\\rho\\rangle/\\rho_{\\rm crit}(t_{\\rm coll})$, and $K/|W|$ is the ratio of total kinetic (i.e. thermal) to gravitational potential energy of the spheres. ", + "conclusions": "" + }, + "0112/astro-ph0112561_arXiv.txt": { + "abstract": "In the Cold Dark Matter (hereafter CDM) scenario even isolated density peaks contain a high fraction of small scale clumps having velocities larger than the average escape velocity from the structure. These clumps populate protoclusters, especially in the peripheral regions, $r \\ge R_{\\rm f}$ (where $ R_{\\rm f}$ is the filtering scale). During the cluster collapse and the subsequent secondary infall, collapsing or infalling clumps (having $v v_{\\rm esc}$. We study the interaction between these two kind of clumps by means of the impulse approximation$^1$ and we find that the collapse of {\\it bound} clumps is accelerated with respect to the homogeneous case (Gunn \\& Gott's model, Ref. 2). The acceleration of the collapse increases with decreasing height of the peak, $\\nu$. We finally compare the acceleration produced by this effect to the slowing down effect produced by the gravitational interaction of the quadrupole moment of the system with the tidal field of the matter of the neighboring proto-clusters studied in Del Popolo \\& Gambera$^3$. We find that the magnitude of the slowing down effect is larger than the acceleration produced by the effect studied in this paper, only in the outskirts of the cluster. We want to stress that the one which we study in this paper is also present in an isolated protocluster, being produced by the interaction of the collapsing clumps with the {\\it unbound} substructure internal to the collapsing clumps itself while that studied in Ref. 3 is produced by substructure external to the density peak. ", + "introduction": "% According to the most promising cosmological scenarios, structure formation is traced back to the evolution of primordial density fluctuations. These fluctuations, originated from quantum fluctuations$^4$$^,$$^5$$ ^,$$^6$$^,$$^7$ in an inflationary phase of the early Universe, grew up through gravitational instability to a maximum radius $r_{\\rm m}$. At the time $t_{\\rm m}$ corresponding to maximum expansion, perturbations broke away from the general expansion and at $\\overline{\\delta} \\sim 1$ them began to collapse. Hence the collapse of perturbations onto local density maxima of the primordial density field has a key role in structure formation and several studies deal with this problem. The problem of collapse has been investigated from two points of view, namely: \\begin{description} \\item i) that of the statistical distribution of the objects formed$^7$$^,$$^8$ \\item ii) that of the structure of these objects and its dependence on the statistical properties of the primordial density field$^2$$^,$$^9$$ ^,$$^{10}$$^,$$^{11}$ $^{12}$$^,$$^{13}$$ ^,$$^{14}$$^,$$^{15}$ $^{16}$$^,$$^{17}$$ ^,$$^{18}$$^,$$^{19}$ $^{20}$$^,$$^{21}$$ ^,$$^{22}$$^,$$^{23}$ $^{24}$. \\end{description} In the spherical accretion model introduced by Hoyle and Narlikar$^{25}$ and applied to clusters of galaxies by Gunn \\& Gott$^{2}$, the matter around the core of the perturbation is a homogeneous fluid with zero pressure. If the density inside the perturbation is greater than the critical density, it is bound and shall expand to a maximum radius $ r_{\\rm m}$: \\begin{equation} r_{\\rm m} = \\frac{r_i}{\\overline {\\delta}} \\end{equation} where $r_{\\rm i}$ is the initial radius and $ \\overline{\\delta}$ is the overdensity inside the radius $ r$. Such matter shall collapse in a time \\begin{equation} T_{\\rm c0}/2 =\\frac{ \\pi}{ H_{\\rm i}} \\frac{(1 +\\overline{\\delta})} {\\overline{\\delta}^{3/2}} \\end{equation} where $ H_{\\rm i} $ is the Hubble parameter at the initial time $ t_{\\rm i} $. \\\\ This model was introduced in order to overcome the problem of the excessively steep density profiles, $\\rho\\propto r^{-4}$, obtained in numerical experiments of simple gravitational collapse. Some authors were able to produce shallower profiles$^2$$^,$$^{26}$$^,$$^{27}$, $\\rho\\propto r^{-2}$, through the ${\\it secondary}$ ${\\it infall}$ process. Several years later, observational evidences for secondary infall in the outskirts of clusters of galaxies has been reported in Ref. 28, 29. \\\\ Even if this model and in general the original form of SIM (secondary infall model) is able to explain better than previous models the structure of clusters of galaxies it has some drawbacks, (e.g. it does not predict the right structure of density profiles), that can be overcome with some improvements in the original model$^{30}$. Two noteworthy limits by the Gunn \\& Gott$^{2}$ model are the following: \\begin{description} \\item {\\it a}) It neglects tidal interaction of the perturbation with the neighboring perturbation; \\item {\\it b}) It neglects substructure existing in the outskirts of proto-galaxies or proto-clusters and in the background. \\end{description} The effect of the tidal field was studied in a recent paper$^{3}$. In that paper, we showed that the gravitational interaction of the quadrupole momentum of a proto-structure with the tidal field of the neighboring proto-structures introduces another potential energy term in the equation describing the collapse which acts in the sense to delay the collapse. This effect has revealed as one of noteworthy importance in solving several of the drawbacks of the CDM model$^{3}$$^,$$^{31}$$^,$$^{32}$$^,$$^{33}$, namely: \\begin{description} \\item 1) lack of a mechanism originating the bias$^{3}$; \\item 2) discrepancies between: \\begin{itemize} \\item a1) X-ray temperature function of clusters and observations$^{31}$. \\item a2) two-point correlation function of clusters and observations$^{31}$; \\item a3) mass function and velocity dispersion function of clusters and observations$^{33}$; \\item a4) shape of clusters and observations$^{34}$; \\item a5) angular two-point correlation function of galaxies and observations$^{32}$. \\end{itemize} \\end{description} For what concerns point \" $b)$\", we may say that during the last two decades, astronomers have discovered that at least a third of galaxy clusters are not dynamically relaxed systems but contain substructure on scales of the same order as the cluster itself$^{34}$$^,$$^{35}$$^,$$^{36}$. This implies that clusters are currently forming or have formed recently enough that they have not had time to undergo significant degree of violent relaxation and phase mixing. When one studies the structure of the velocity field around a peak one finds that at radii $ r$ larger than the filtering scale $ R_{\\rm f}$, ($r \\ge R_{\\rm f}$), a sensitive fraction of the matter of the peak is gravitationally unbound to it$^{37}$. This is not surprising since a non zero fraction of unbound matter should be expected even in the case of a uniform system with a Maxwellian velocity distribution (see section 2.2 for a discussion). The encounters between collapsing clumps and the unbound high velocity clumps influence the collapse of the proto-cluster. We want to study this point and to compare its effect on cluster collapse with that produced by the effects of tidal interaction of the protocluster with the neighboring ones. \\\\ In the following we show how the unbound substructures influences the collapse of the perturbation. The plan of this work is the following: in Section ~2 we deal with the unbound substructure effects. In Section ~3 we give our results and finally Section ~ 4 is devoted to the conclusions. ", + "conclusions": "In this paper we have studied the role of unbound high velocity clumps on the collapse of density peaks in a SCDM model ($\\Omega_0=1$, $h=0.5$, $n=1$). We have shown that: \\\\ 1) the encounters of collapsing shells with unbound clumps produce an acceleration in the collapse of density peaks. The effect is larger for peaks having a lower value of $ \\nu$ (see Fig.~2 $ \\div $ 3). \\\\ 2) A comparison between the collapse acceleration due to encounters of shells with unbound clumps and the slowing down effect produced by tidal interaction between a protocluster and the neighboring ones (see Ref. 3) has shown that these two effects act in opposite direction, namely the first accelerate the shell collapse, while the second delay the protocluster collapse, especially in its outskirts. The magnitude of the two effects is comparable at $1 h^{-1} {\\rm Mpc}$ but the effect of the tidal field dominates at larger radii. \\\\ 3) The acceleration in the collapse is due to the additive action of several high speed clumps with the matter of the infalling shells." + }, + "0112/astro-ph0112082_arXiv.txt": { + "abstract": "This paper investigates the detailed dynamical properties of a relatively homogeneous sample of disc-dominated S0 galaxies, with a view to understanding their formation, evolution and structure. By using high signal-to-noise ratio long-slit spectra of edge-on systems, we have been able to reconstruct the complete line-of-sight velocity distributions of stars along the galaxies' major axes. From these data, we have derived both model distribution functions (the phase density of their stars) and the approximate form of their gravitational potentials. The derived distribution functions are all consistent with these galaxies being simple disc systems, with no evidence for a complex formation history. Essentially no correlation is found between the characteristic mass scale-lengths and the photometric scale-lengths in these galaxies, suggesting that they are dark-matter dominated even in their inner parts. Similarly, no correlation is found between the mass scale-lengths and asymptotic rotation speed, implying a wide range of dark matter halo properties. By comparing their asymptotic rotation speeds with their absolute magnitudes, we find that these S0 galaxies are systematically offset from the Tully-Fisher relation for later-type galaxies. The offset in luminosity is what one would expect if star formation had been suddenly switched off a few Gyrs ago, consistent with a simple picture in which these S0s were created from ordinary later-type spirals which were stripped of their star-forming ISM when they encountered a dense cluster environment. ", + "introduction": "Since S0 galaxies seem to have properties intermediate between elliptical and spiral galaxies, Hubble (1936) placed these gas-poor disc systems between the spirals and ellipticals, at the bifurcation of his tuning fork galaxy classification scheme. However, three-quarters of a century later, it remains an open question as to whether this arrangement in any way reflects the physical origins of S0 galaxies. One approach to addressing question of this kind has been to search directly for signs of evolution in galaxy morphologies by observing samples of these systems over a wide range of redshifts. In a classic study of this kind, Butcher \\& Oemler (1978) found that the fraction of blue galaxies in clusters was much higher in the past than it is now. They surmised that blue spirals were being converted to earlier-type systems. This discovery fitted in extremely well with Gunn \\& Gott's (1972) suggestion that S0 galaxies could form from spiral galaxies in the dense environment of a cluster via tidal and ram pressure stripping. More recently, high-resolution observations (e.g. Couch et al. 1994; Lavery, Pierce \\& McClure 1992) have confirmed that the blue galaxies in high redshift clusters are, indeed, relatively normal spiral galaxies, which could well be the progenitors of S0 systems. The only problem with such analyses is that one is ne\\-ces\\-sa\\-rily observing different galaxies in the nearby and distant samples, so any inference about the evolution from one type to another is of a circumstantial nature. A convincing case therefore requires that one look in some detail at the ``finished articles,'' in order to find any archaeological evidence for the proposed evolution. Since galaxies are intrinsically dynamical entities, one should be able to extract important clues from their stellar kinematics, as derived from absorption-line spectra, as well as their photometric pro\\-per\\-ties. With high quality spectral data, one can measure the full line-of-sight velocity distribution (e.g. Koprolin \\& Zeilinger 2000, Fisher 1997), providing information not only on the motions of the stars, but also the gravitational potential responsible for the motions. For S0 galaxies, one important piece of stellar-dynamical evidence would be provided by an analysis of the Tully-Fisher relation. In later-type disc galaxies, there is a strong correlation between circular rotation speed and optical luminosity (Tully \\& Fisher 1977), and this relationship becomes even tighter when one looks in the near infrared (Pierce \\& Tully 1992). If S0 galaxies formed in a relatively benign way from spiral galaxies, one would expect that their optical luminosities and circular rotation speeds would be little affected by the process, so a Tully-Fisher relation should still be apparent; the only significant difference would be that the stripping of their gas would switch off the star formation process, so the optical luminosities of the S0s should fade over time, shifting the zero-point of the relation. To-date, the most thorough search for a Tully-Fisher relation in S0 galaxies was that made by Neistein {\\it et al.}\\ (1999). This analysis is complicated by the fact that there is no simple measure of the circular speed in an S0 galaxy: there is no gas at large radii moving on circular orbits, and the stars follow significantly elliptical orbits, so their mean streaming velocity at any radius is lower than the local circular velocity. Neistein {\\it et al.} overcame the problem by appealing to the equations of galactic dynamics: by making a few simplifying assumptions, they were able to use the asymmetric drift equation (Binney \\& Tremaine 1987) to combine the mean streaming velocity and velocity dispersion of the stars in order to estimate the circular velocity at each radius. Although this analysis revealed some evidence for a trend between I-band luminosity and circular speed, Neisten {\\it et al.} found a huge scatter in the relation, and that there was very little offset in the mean from the Tully-Fisher relation of later-type galaxies. They therefore concluded that these systems could not all have formed from the simple stripping of spiral galaxies. Instead, they suggested that the S0 classification actually represents a rather heterogeneous class of galaxies, which formed through a rather wide variety of processes. They further suggested that the absence of an offset in the Tully-Fisher relation could be understood if the S0 galaxies have more massive discs, so any fading in their luminosities is offset by the larger number of stars. In order to explore these conclusions a little further, this paper presents a detailed dynamical analysis of six S0 galaxies. These galaxies have been selected to contain relatively small central bulges. If any S0s formed from simple stripping processes, one would expect their bulges to be little affected. Thus, by selecting S0s with the small bulges characteristic of later-type spirals, one might hope to pick out a relatively homogeneous subsample of systems that formed via this route. Unfortunately, it is only when S0 galaxies are very close to edge on that one can reliably determine that their bulges are small. As Neistein {\\it et al.}\\ (1999) demonstrated, the line-of-sight integration of starlight through such an edge-on disc means that there is quite a large correction to convert the observable mean line-of-sight velocity into the circular streaming velocity of the stars. As a further complication, the observed velocity distribution for any line of sight through an edge-on disc will be highly non-Gaussian, due to the contribution from stars at large radii with small line-of-sight velocities. A dynamical analysis based on moments derived from a Gaussian fit to the line-of-sight velocity distribution is therefore prone to systematic error. To obviate these difficulties, the current analysis uses data with a high enough signal-to-noise ratio for the complete line-of-sight velocity distribution to be derived, and these data are then fitted to a complete dynamical model in order to derive both the stellar distribution function and the gravitational potential needed for the Tully-Fisher relation. The remainder of the paper is laid out as follows. Section~\\ref{datasec} describes the sample, the data analysis, the model fitting, and the resulting distribution functions and rotation curves. Section~\\ref{discussionsec} discusses the correlations in derived quantities, with particular reference to the Tully-Fisher relation. Section ~\\ref{conclusionsec} presents the conclusions that can be drawn from this analysis. ", + "conclusions": "In this paper, we have calculated the first detailed dynamical models for a small sample of edge-on S0 galaxies with small bulges. In addition to producing distribution functions for these disc-dominated systems, which look very much as one would expect for normal disc systems, the analysis also returned estimates for the parameters of their gravitational potentials. The interpretation of these data all points to a simple picture in which these systems were formed by the stripping of gas from normal spiral galaxies. The distribution functions are all well modeled by unexceptional stellar discs, similar to those expected in the old stellar populations of spiral galaxies. In addition, the galaxies obey a reasonably tight Tully-Fisher relation, which is offset from the relation for normal spiral galaxies by the amount that one would expect if star formation had been shut off a few Gyrs ago, so that all that remains in these systems are the rather fainter old stellar populations. This result appears to conflict with Neistein {\\it et al.}'s (1999) analysis, which showed a much greater scatter in the Tully-Fisher relation with less systematic offset. Part of the difference may be due to the lower signal-to-noise ratio of the data in their larger sample, which limited their ability to carry out detailed dynamical modeling, particularly for galaxies that lie very close to edge-on. However, there is also a systematic difference in the way that the samples were selected: the edge-on galaxies in the analysis of this paper were specifically chosen to contain small bulges. This selection criterion means that these galaxies are prime candidates to have formed from gas-stripped spiral galaxies. If, as Neistein {\\it et al.} suggest, S0s are a ``mixed bag'' that formed in a variety of ways, it should come as no surprise that this particular subsample obey a tight Tully-Fisher relation that is not seen in the general population of S0s. To test such hypotheses, we ultimately need data of the quality presented in this paper for a much larger sample of galaxies, extending both the range of absolute magnitudes and of other parameters such as the bulge size." + }, + "0112/astro-ph0112031_arXiv.txt": { + "abstract": "I review the main characteristics of structure formation in the quintessential Universe. Assuming equation of state $w=p/\\varrho=$const I provide a brief description of the background cosmology and discuss the linear growth of density perturbations, the strongly nonlinear evolution, the power spectra and rms fluctuations as well as mass functions focusing on the three values $w=-1, -2/3$ and $-1/3$. Finally I describe the presently available and future constraints on $w$. ", + "introduction": "Our knowledge of background cosmology has recently improved dramatically due to new supernovae and cosmic microwave background data. Current observations favor a flat Universe with matter density $\\Omega_0=0.3$ \\cite{hr} and the remaining contribution in the form of cosmological constant \\cite{cpt, llpr} or some other form of dark energy. The models with cosmological constant are known, however, to suffer from two major problems. One is related to the origin of the constant - it cannot be explained in terms of the vacuum energy since its energy is orders of magnitude smaller. The other is the lack of explanation why the present densities in matter and cosmological constant are comparable. A new class of models that solve these problems and also satisfy present observational constraints has been proposed a few years ago \\cite{cds}. In these models the cosmological constant is replaced with a new energy component, called quintessence, characterized by the equation of state $p/\\varrho=w \\neq -1$. The component can cluster on largest scales and therefore affect the mass power spectrum \\cite{ma} and microwave background anisotropies \\cite{balbi, bacci}. The investigations of the physical basis for the existence of such component are now more than a decade old \\cite{rp}. One of the promising models is based on so-called ``tracker fields\" that display an attractor-like behavior causing the energy density of quintessence to follow the radiation density in the radiation dominated era but dominate over matter density after matter-radiation equality \\cite{zws, swz}. It is still debated, however, how $w$ should depend on time, and whether its redshift dependence can be reliably determined observationally \\cite{bm, mbs, wa}. A considerable effort has gone into attempts to put constraints on models with quintessence and presently the values of $-1 5-10$~arcmin). Many of the fields however suffer severely from the grating response and from confusing thermal emission from nearby strong sources. \\citet{YoursTruely2000} used the GMRT \\citep{GMRT} at 327~MHz to confirm three SNRs selected from the survey done by \\citet{DUNCAN_SNRFrom2.4GHzSurvey} at 2.4 GHz using the Parkes 64-m dish. The GMRT at 327~MHz provides a resolution of $\\sim20$~arcsec for observations in the southern sky and is sensitive to spatial scales of up to 30~arcmin. At this frequency, thermal emission from typical HII regions is weak while the emission from SNRs remain relatively strong. The relatively smaller field-of-view ($\\sim 1^\\circ.4$ at 327~MHz) of the GMRT offers a further advantage in terms of attenuating confusing emission farther away. A relatively better imaging performance of the GMRT compared to that of the MOST allows higher dynamic range mapping of complicated fields like the ones imaged by \\citet{Gray-III} and hence more reliable maps at a higher resolution. Thus the combination of high angular resolution, sensitivity to large scale structures and comparatively higher sensitivity that the GMRT provides at these low frequencies makes it a good instrument for studies of Galactic SNRs. I present here the results of 327~MHz GMRT observations of three objects (namely, G001.4$-$0.1, G003.8$+$0.3, and G356.3$-$1.5) which have been included in the Galactic SNR Catalogue \\citep{SNRCAT} based on 843-MHz \\citep{Gray-III} observations alone, and one candidate SNR (G004.2$+$0.0). These fields were selected from the 843-MHz survey by \\citet{Gray-III}. ", + "conclusions": "In conclusion, we confirm G001.4$-$0.1, G003.8$+$0.3 and G356.3$-$1.5 to be Galactic SNRs based on the morphology and non-thermal nature of emission from these objects. These SNRs were reported in the list of candidate SNRs by \\citet{Gray-III} and have since then been include in the latest catalogue of Galactic SNRs \\citep{SNRCAT} based on the morphological evidence alone. These observations add the new information of the non-thermal nature of emission from these objects which firmly establishes them as Galactic SNRs. Significant thermal emission is seen at the location of another small diameter candidate SNR G004.2$+$0.0 reported by \\citet{Gray-III}. The radio flux densities of this object are consistent with it being a thermal source. Morphology of G001.4$-$0.1, coincidence of a compact \\OHMASER\\ spot with the western arc and morphologically similar extended \\OHMASER\\ emission towards this sources is suggestive of an interaction with a nearby molecular cloud." + }, + "0112/astro-ph0112177_arXiv.txt": { + "abstract": "We present the results of a detailed timing analysis of observations of Cen~X-3 taken by the University of Durham Mark-6 imaging atmospheric Cherenkov telescope in 1997-1999. The presence of a TeV $\\gamma$-ray signal at the overall $\\geq 4.5\\sigma$ significance level in the `fully cut' image selected data, as reported earlier, is confirmed. A search for possible modulations of $\\gamma$-rays with the pulsar spin period $P_{0}\\approx 4.8 \\,\\rm s$ was performed by the step-by-step application of image parameter cuts of gradually increasing hardness. The data of each of 23 days of observations have not revealed any statistically significant Rayleigh power peak, except for 1 day when a peak with a chance probability $p=6.8\\times 10^{-7}$, was found in `soft-cut' data sets. This modulation, if real, is blue shifted by 6.6 msec ($>10^3$ kms$^{-1}$) from the nominal second harmonic of the X-ray pulsar. Taking the large number of frequency trials into account, the estimated final probability of such a peak by chance still remains $< 10^{-2}$. Bayesian statistical analysis also indicates the presence of such modulations. We show that the behaviour of the Rayleigh peak disappearing in the fully cut data set is actually quite consistent with the hypothesis of a $\\gamma$-ray origin of that peak. No modulation of the VHE $\\gamma$-ray signal with the pulsar orbital phase is found. In the second part of the paper we consider different theoretical models that could self-consistently explain the existing data from Cen~X-3 in high-energy (HE, $E\\geq 100\\,\\rm MeV$) and very high energy (VHE, $E\\geq 100\\,\\rm GeV$) $\\gamma$-rays. We propose on the basis of the energetics required that all reasonable options for the $\\gamma$-ray production in Cen~X-3 must be connected to jets emerging from the inner accretion disc around the neutron star. One of the principal options is a large-scale source, with $R_{\\rm s} \\sim 10^{13} - 10^{14} \\,\\rm cm$\\,; this assumes effective acceleration of electrons up to $\\sim 10 \\,\\rm TeV$ by shocks produced by interaction of these jets with the dense atmosphere of the binary. It is shown that such a quasi-stationary model could explain the bulk of the $\\gamma$-radiation features observed except for the $\\gamma$-ray modulations with the pulsar spin. These modulations, if genuine, would require an alternative source with $R_{\\rm s} \\ll 10^{11}\\,\\rm cm$. We consider two principal models, hadronic and leptonic, for the formation of such a compact source in the jet. Both models predict that the episodes of pulsed $\\gamma$-ray emission may be rather rare, with a typical duration not exceeding a few hours, and that generally the frequency of pulsations should be significantly shifted from the nominal frequency of the X-ray pulsar. The opportunities to distinguish between different models by means of future $\\gamma$-ray observations of this X-ray binary are also discussed. ", + "introduction": "Centaurus X-3 has been one of the prominent galactic sources of hard radiation since its discovery as one of the first cosmic X-ray sources (Chodil et al. \\cite{Chodil67}), and was the first X-ray pulsar to be discovered in a binary system (Giacconi et al \\cite{G71}, Schreier et al. \\cite{Schreier72}). All the basic parameters of this archetypal high mass X-ray binary are well known. The pulsar has a spin period $P_0 \\approx 4.8\\,\\rm s$ and an orbital period $P_{\\rm orb} \\approx 2.1 \\,\\rm d $, with a gradual shortening (i.e. `spinning-up') of both periods in time, and with a deep eclipse of the X-ray source at the orbital phases $-0.12 \\leq \\phi \\leq 0.12$ (for reviews see Joss \\& Rappaport \\cite{Joss84}, Nagase \\cite{Nagase89}). The X-ray luminosity of the pulsar is very large, reaching $L_{\\rm X} \\sim 10^{38} \\,\\rm erg s^{-1}$ in the `high' state (e.g. White et al. \\cite{White83}, Burderi et al. \\cite{Burderi00}), which implies a massive accretion of material onto the neutron star from the optical companion (Pringle \\& Rees \\cite{Pringle72}, Lamb et al. \\cite{Lamb73}). The optical companion (V779 Cen) was discovered by Krzeminski (\\cite{Krz74}) and has been identified as an evolved O-type star with surface temperature $T\\geq 3\\times 10^4 \\,\\rm K$ at a distance $\\sim 8\\,\\rm kpc$ from the Sun, and a bolometric magnitude of $M_{\\rm bol} \\sim -9$ (Hutchings et al. \\cite{Hut79}). The masses of the stars in this binary are estimated as $M_{\\rm n}\\simeq 1.2 \\, M_{\\odot}$ and $M_{\\rm O-star} \\simeq 20\\,\\rm M_{\\odot}$. The value of $\\rm P_{orb}$ suggests the separation between the centres of the stars to be $a=1.3\\times 10^{12}\\, \\rm cm$ and the radius of the massive star filling its Roche lobe $R_{\\rm O} = 8.6\\times 10^{11}\\,\\rm cm$ (see e.g. Clark et al. \\cite{Clark88}, Ash et al. \\cite{Ash99}). Cen X-3 is also known as one of those X-ray binaries from which $\\gamma$-ray signals have been reported. In the domain of high energy (HE) $\\gamma$-rays, conventionally $E\\geq 100 \\,\\rm MeV$, the $\\gamma$-ray flux during 2 weeks of observations of Cen X-3 in October 1994 by EGRET was at the level $F(>100\\,\\rm MeV) =(9.2\\pm 2.3)\\times 10^{-7} \\,\\rm ph\\, cm^{-2} s^{-1}$ (Vestrand et al. \\cite{V97}). This excess flux was significant at a $5\\sigma$ level, which is generally considered as a reliable detection with the EGRET instrument. The timing analysis has shown a significant modulation of the signal with the pulsar spin, precisely at the contemporaneous frequency of X-ray pulsations measured by BATSE, with a chance probability estimate based on the Rayleigh test statistic of about $1.6\\times 10^{-3}$. No modulation with the orbital phase of the pulsar was found: the $\\gamma$-ray signal seems to be quite homogeneously distributed throughout the entire orbit, with 68 of the 264 HE $\\gamma$-rays from Cen X-3 being detected in the pulsar eclipse orbital phase $| \\phi | \\leq 0.12$. This suggests a production site of the $\\gamma$-rays far away from the pulsar. The $\\gamma$-ray signal was not found in the data of other observing periods of Cen X-3 by EGRET, which suggests a significant variability of the HE $\\gamma$-ray source on a time scale of several months (see Vestrand et al. 1997 for details). In the domain of very high energies (VHE), conventionally $E\\geq 100\\,\\rm GeV$, sporadic $\\gamma$-ray signals from Cen X-3 had been earlier reported by the University of Durham (Brazier et al. \\cite{Bra90a}) and the Potchefstroom groups (North et al. \\cite{N90}). Note that the early detections of VHE $\\gamma$-rays from X-ray binaries in the 80's had been carried out by non-imaging Cherenkov telescopes which had rather poor sensitivity. The $\\gamma$-ray signals were therefore extracted mostly on the basis of the timing analyses, and are subject to controversy concerning their reliability (see Weekes \\cite{Weekes92}). It is worth noting in this regard that Cen X-3 is the first and until now the only X-ray binary to be detected as a source of $E \\geq 400\\,\\rm GeV$ $\\gamma$-rays by a contemporary imaging instrument, the Durham Mark-6 telescope, as we have reported earlier (Chadwick et al. \\cite{C98,C00}). The excess ($\\gamma$-ray) signal has been found in the `ON-source' data during each of the 3 years of observations from 1997 to 1999. The estimated mean $\\gamma$-ray flux is about $F(>400\\,\\rm GeV) \\simeq 2.8 \\times 10^{-11} \\,\\rm cm^{-2} s^{-1}$, at the significance level for the entire data set of $4.7\\sigma$. No significant modulations in the combined ON-source data with either the 4.8\\,s pulsar period or 2.1\\,d orbital period of the binary were found, but these data were analyzed after application of the image cut procedures (Chadwick et al. \\cite{C00}) which we shall see later can be very counter productive in a periodicity search. In this paper we present the results of a more detailed timing analysis of observations of Cen X-3, which in particular includes a step-by-step search for possible short-term episodes of $\\gamma$-ray emission modulated with the spin of the pulsar, as well as a search for a modulation of the $\\gamma$-ray signal with the pulsar orbital phase (Section 2). Then we carry out in Section 3 a detailed theoretical study of the consistency of different models for the production of $\\gamma$-rays in Cen X-3 with the experimental data currently available in both HE and VHE domains. Finally, in Section 4 we summarize the results of our current study, and discuss some possible tests for future $\\gamma$-ray observations that could help to distinguish between different models of $\\gamma$-ray production in this X-ray binary. ", + "conclusions": "The analysis of the data of observations of Cen X-3 accumulated during 23 days of observations of Cen X-3 with the University of Durham Mark-6 imaging telescope during 1997-1999 confirms the presence of a $\\gamma$-ray signal at the overall significance level $\\geq 4.5\\,\\sigma$ in agreement with our previous results (Chadwick et al. \\cite{C00}). The behaviour of the signal, which does not drop but rather increases for images with enhanced brightness (see Table 1), indicates that the spectra of VHE $\\gamma$-rays from Cen X-3 are probably much harder than the spectra of ambient cosmic rays with $\\alpha_{\\rm CR} \\approx 2.7$. No indications for any correlation of the VHE signal with the orbital motion of the pulsar, including the pulsar eclipse phase, is found. Given that a similar result is found by EGRET for HE $\\gamma$-rays detected from Cen X-3, this effect excludes any close vicinity of the pulsar as a possible site for $\\gamma$-ray production. The analysis of the shower arrival times does not reveal any statistically significant peak of Rayleigh powers in most of the data, except for 1 observation day, the 21 Feb 1999, showing a strong peak with an estimated probability of occurence of such a peak by chance, after taking into account the large number of IFF trials, of significantly below $10^{-2}$. The analysis using Bayesian statistics results in an even lower final probability, $p_{\\rm Bay} \\simeq 7.5\\times 10^{-4}$. It is indicative that the peak powers in both statistics are found in the data after application of soft image cuts, whereas they completely disappear when the full (hard) image cuts maximizing significance in the overall data of observations are applied. This behaviour is just what one should expect assuming that the modulated signal in the data of 21 Feb 1999 is really connected with VHE $\\gamma$-rays. The pulsed signal significantly increases with the threshold energy of the events, which indicates that the spectrum of pulsed events is hard as well. The position of modulations found in 21 February 1999 data is blueshifted from the nominal {half period of the X-ray pulsations} by $\\delta \\nu /\\nu \\simeq 3.7 \\times 10^{-3}$. Such shifts should be reasonably expected in any model where the $\\gamma$-rays are produced far away from the pulsar. Moreover, on the basis of the theoretical modelling in this paper, which suggests high speeds of compact source(s) needed for production of pulsed VHE $\\gamma$-ray emission, we would generally expect that $\\gamma$-ray pulsations could show even larger shifts, up to $|\\Delta \\nu | /\\nu_0 \\sim 0.3$ in the case of ejecta moving with the speed $\\leq 0.3 c$ such as in SS~433. The principal models for $\\gamma$-ray production in Cen X-3 should be able to produce VHE radiation at distances significantly beyond the orbit of the pulsar, effectively at $R\\geq 3\\times 10^{12}\\,\\rm cm$. Meanwhile, high energies needed for the $\\gamma$-ray fluxes in this binary can be provided only by the gravitational potential of the neutron star. Confirmation of the $\\gamma$-ray fluxes from Cen X-3 by future observations would therefore confirm production of powerful jets by the inner accretion disc of this X-ray binary. Except for the phenomenon of $\\gamma$-ray pulsations (which obviously needs confirmation by forthcoming $\\gamma$-ray detectors), an acceptable interpretation of the average fluxes currently reported in HE and VHE $\\gamma$-ray regions can in principle be provided in the framework of both spatially extended and compact source models, which may also `co-exist' operating together. The principal difference between these 2 model approaches is that the extended $\\gamma$-ray source implies quasi-stationary emission on times scale of at least several days, and probably even weeks, whereas both hadronic and leptonic compact source models predict fast evolution of the $\\gamma$-ray fluxes on time scales of a few hours, and with the `disappearance' of the source in a few days. Another informative feature could be the behaviour of the spectral fluxes at both high and very high energies. An important prediction of both compact source models is that pulsed $\\gamma$-ray emission can {\\it only} be episodic, with a typical duration of no more than a few hours. For $\\gamma$-rays of $E \\sim 100 \\,\\rm GeV$ which are to be produced at large distances this is practically a model-independent requirement. However, $\\gamma$-rays with $E\\leq 10\\,\\rm GeV$ can escape from the binary, even being produced relatively (but not very!) close to the X-ray pulsar. At these energies, therefore, detection of pulsed $\\gamma$-rays in the eclipse phase of the X-ray pulsar would be very informative. \\vspace{1mm} We can expect that future observations of Cen X-3 and other X-ray binaries with forthcoming sensitive $\\gamma$-ray detectors in the $\\geq 50 \\,\\rm MeV$ (GLAST) and $\\geq 50\\,\\rm GeV$ (HESS, CANGAROO-III, VERITAS) domains will provide large amounts of key information about high energy processes in these powerful galactic objects." + }, + "0112/astro-ph0112494_arXiv.txt": { + "abstract": "Silicate dust grains in the interstellar medium are known to be mostly amorphous, yet crystalline silicate grains have been observed in many long-period comets and in protoplanetary disks. Annealing of amorphous silicate grains into crystalline grains requires temperatures $\\gtsimeq 1000 \\, {\\rm K}$, but exposure of dust grains in comets to such high temperatures is incompatible with the generally low temperatures experienced by comets. This has led to the proposal of models in which dust grains were thermally processed near the protoSun, then underwent considerable radial transport until they reached the gas giant planet region where the long-period comets originated. We hypothesize instead that silicate dust grains were annealed {\\it in situ}, by shock waves triggered by gravitational instabilities. We assume a shock speed of $5 \\, {\\rm km} \\, {\\rm s}^{-1}$, a plausible value for shocks driven by gravitational instabilities. We calculate the peak temperatures of micron and submicron amorphous pyroxene grains of chondritic composition under conditions typical in protoplanetary disks at 5 -- 10 AU. Our results also apply to chondritic amorphous olivine grains. We show that {\\it in situ} thermal annealing of submicron and micron-sized silicate dust grains can occur, obviating the need for large-scale radial transport. ", + "introduction": "Comets are collections of the dust grains and ices that resided in the solar nebula. The dust grains and ices in comets do not appear to have been subject to any processing mechanisms since the comets formed at the birth of our solar system (Hanner 1999; Brucato et al.\\ 1999). Cometary material therefore is thought to contain pristine solar nebula material and to represent the physical conditions of the solar nebula at the birth of the comet, including temperatures $\\ltsimeq 50 \\, {\\rm K}$ (Meier \\& Owen 1999). Complicating this interpretation is the spectroscopic evidence of silicate grains processed at temperatures $\\gtsimeq 1000 \\, {\\rm K}$ existing in comets. Crystalline silicates with distinct mid- and far-IR spectral resonance features can form a significant component of dust (30 -- 50\\%) in comet comae (Wooden et al.\\ 1999; Harker et al.\\ 2002). The majority of the crystalline silicate grains seen in many comets must have been condensed or annealed (presumably by thermal processing) in the solar nebula. As pointed out by Irvine et al.\\ (2000) and Wooden et al.\\ (1999), despite the observations of crystalline olivines [$({\\rm Mg}_{x}{\\rm Fe}_{x-1})_{2}{\\rm SiO}_{4}$] and pyroxenes [$({\\rm Mg}_{x}{\\rm Fe}_{x-1}){\\rm SiO}_{3})$] ($x$ ranges between 1 [Mg-pure] and 0 [Fe-Pure]) in circumstellar environments (Waters \\& Waelkens 1998; Malfait et al.\\ 1998), crystalline silicates are essentially unseen in the interstellar medium ($< 5$\\%; Li \\& Draine 2001). However, it is not clear how silicate grains could be heated to $\\sim 1000 \\, {\\rm K}$ in the cold, comet-forming environment ($\\sim 5 - 100 \\, {\\rm AU}$). In the solar nebula, the high temperatures $\\gtsimeq 1000 \\, {\\rm K}$ needed to anneal amorphous interstellar silicate grains, or to condense crystalline silicate grains directly from the gas phase ($\\gtsimeq 1300 \\, {\\rm K}$) were achieved in steady-state only at distances $\\ltsimeq 1 \\, {\\rm AU}$ from the protoSun at early times (e.g., Bell et al.\\ 1997). Proposed mechanisms for transporting crystalline silicates to the comet forming zones include radial diffusion from vigorous thermal convection (Nuth 2001; Hill et al.\\ 2001), or outflows driven by reconnecting magnetic field lines (``X-winds'' [Shu et al.\\ 1996]). Although X-winds have been invoked as the cause of chondrule formation (mm-sized spheres of silicate rock found in meteorites [see Jones et al.\\ 2000]), there is compelling evidence (chondrule rim thickness, chondrule-matrix complementarity) that chondrules formed {\\it in situ}, in the asteroid belt region (see summary in Desch \\& Connolly 2002; hereafter DC02), obviating the need for radial transport. The similarities between the annealing of silicate grains and the melting of chondrules have been pointed out by Grady et al.\\ (2000) and Nuth (2001). These similarities between chondrule formation and silicate annealing lead us to hypothesize that interstellar, amorphous silicate grains were also annealed {\\it in situ}, at 5 -- 10 AU, by shock waves in the solar nebula. In this {\\it Letter}, we calculate the thermal histories of submicron silicate grains passing through nebular shocks, using the code of DC02 that was used to calculate the thermal histories of chondrules. We conclude that silicate grains were likely annealed {\\it in situ} in the gas giant planet region by nebular shocks, and we attribute these shocks to disk gravitational instabilities. ", + "conclusions": "Annealing of silicate grains can be accomplished by shocks in the outer solar nebula (at 10 AU). Smaller grains are more easily annealed due to their poorer radiative efficiencies. These results are not affected by the presence of ice mantles, because such mantles would sublimate before reaching the shock front. Water ice is at least as good an absorber as silicates of the thermal radiation emanating from the shock front, and dust grains can easily absorb enough energy from the radiation field to heat above 150 K. Figure~1$a$ shows that a dust grain attains temperatures of $> 1000$~K for $> 10$~minutes. Since the grain is in radiative equilibrium, the grain is absorbing and emitting thermal radiation at a rate of $4\\pi a_{\\rm p}^{2} \\, Q_{\\rm p} \\, \\sigma T_{\\rm p}^{4}$. If the grain is assumed to be made entirely of water ice with latent heat of sublimation $2.83 \\times 10^{10}$~erg~g$^{-1}$ and radiative absorption efficiency $Q = 0.1$, it can easily be shown that the ice will completely evaporate in $< 0.2$~s. This evaporation time is much less than the 10~minutes the particle stays at temperatures $> 1000$~K. Shocks raise grains to high temperatures whether or not they have ice mantles. Therefore, provided shock speeds $\\sim 5 \\, {\\rm km} \\, {\\rm s}^{-1}$ can be attained, annealing of silicates {\\it in situ} is very likely. Shocks are very likely to occur early in the evolution of the solar nebula. The accretion disk models of Bell et al.\\ (1997) predict that T Tauri disks with mass accretion rates $10^{-8} \\, M_{\\odot} \\, {\\rm yr}^{-1}$ and $\\alpha = 10^{-2}$ are locally gravitationally unstable beyond 11 AU, and disks with $\\alpha = 10^{-4}$ are locally unstable beyond 6 AU. Since mass accretion rates $\\sim 10^{-8} \\, M_{\\odot} \\, {\\rm yr}^{-1}$ are likely to persist for many Myr in T Tauri systems (Gullbring et al.\\ 1998), gravitational instabilities are likely to persist in the outer regions of protoplanetary disks. Since the gravitational instabilities manifest themselves as global modes (Laughlin \\& Rozyczka 1996), they are likely to affect the bulk of the disk, even where it is locally stable. How the gravitational instabilities in a marginally unstable disk manifest themselves has been explored numerically by Boss (2000, 2001a,b), who discovered the development of gravitationally bound clumps and highly non-axisymmetric density structures associated with the clumps in such disks. These structures would be overdense by orders of magnitude, oriented more or less radially, extending inward from the clumps and co-rotating with them (see Figure 1 of Boss [2000]). Gas orbiting closer to the Sun would orbit faster than these density structures and collide more or less normally with them, at relative velocities comparable to the difference in Keplerian angular velocities between the orbiting gas and the clump driving the structure. For example, a clump at 10 AU would drive a density structure with an orbital period $\\approx 30 \\, {\\rm yr}$; material orbiting every $15 \\, {\\rm yr}$ at 6 AU would slam into this structure at a relative velocity $\\sim 6 \\, {\\rm km} \\, {\\rm s}^{-1}$. If nebular shocks existed in the outer solar nebula, whatever their source, we find that submicron silicate grains could readily be annealed by $5 \\, {\\rm km} \\, {\\rm s}^{-1}$ shocks, if the gas densities were $\\gtsimeq 10^{-10} \\, {\\rm g} \\, {\\rm cm}^{-3}$. These densities are plausible for the 5 -- 10~AU region of the solar nebula, according to the $\\alpha$ disk models of Bell et al.\\ (1997), but not beyond 10 AU. Therefore, comets forming in the Kuiper Belt Region (greater than 35~AU) should not contain crystalline silicate dust annealed by shocks. Radial diffusion by turbulence is unlikely to change this result. Within the time a comet takes to form, $t \\sim 10^{5} \\, {\\rm yr}$ (Weidenschilling 1997), turbulent diffusion mixes gas across distances $\\sim (\\alpha \\, C H t)^{1/2}$, where $C$ is the sound speed and $H$ the scale height of the nebula. For $\\alpha \\approx 10^{-3}$, the mixing distance over $10^{5}$ years at 10 AU is $\\ltsimeq 5 \\, {\\rm AU}$; over $10^{6}$ years it is $\\ltsimeq 15 \\, {\\rm AU}$. The vast majority of short-period comets, which come from the Edgeworth-Kuiper belt and which mostly formed beyond 35 AU, should contain amorphous grains only. The few short-period comets with crystalline silicates must have been scattered to the Edgeworth-Kuiper belt after forming at 5 -- 10~AU, which is possible dynamically (Duncan \\& Levison 1997). However, most comets formed at 5 -- 10~AU will instead be scattered into the Oort cloud and be observed entering the solar system today as long-period comets. It is in long-period comets that crystallinity should be expected. The observational data confirm the prediction that crystalline silicates should be found almost exclusively in long-period comets. Fits to Hale-Bopp's (C/1995 O1) 11.2~\\micron\\ features have led to the determinations that, even pre-perihelion, $\\approx 20 \\%$ of all silicates (Hanner et al. 1997, Wooden et al.\\ 1999), and $\\approx 30 \\%$ of the olivine silicates (Harker et al.\\ 2002) must be in crystalline form. According to Yanamadra-Fisher \\& Hanner (1999), comets with comparable peaks at 10 and 11.2~\\micron, such as P/Halley, Bradfield 1987 XXIX, Mueller, Levy 1990 XX, and Hale-Bopp, are well fit by a combination of 20\\% crystalline and 80\\% amorphous silicates. The crystalline silicate feature at $11.2 \\, \\mu{\\rm m}$ has also been observed in Hyakutake (C/1996 B2; Mason et al.\\ 1998), and McNaught-Hartley (C/1999 T1; Lynch, Russell \\& Sitko 2001). Not all long-period comets have crystalline silicates ---Hanner et al.\\ (1994) found no evidence for annealing in Kohoutek 1973 XII, Austin 1990 V, Okazaki-Levy-Rudenko 1989 XIX, and Wilson 1987 VII ---but most do. On the other hand, only in one short-period comet, 103/P Hartley 2, have crystalline silicates been observed (Crovisier et al.\\ 1999). Data are inconclusive or negative in others (Harker et al.\\ 1999; Hanner et al.\\ 1996), including comets P/Brorsen-Metcalf 1989 X (Hanner et al.\\ 1994), P/Churyumov-Gerasimenko (Hanner et al.\\ 1985), P/Grigg-Skjellerup (Hanner et al.\\ 1984), and P/Schaumasse (Hanner et al.\\ 1996), each indistinguishable from blackbodies (Hanner et al.\\ 1994). Winds hypothetically carrying annealed grains from the inner solar system would preferentially seed Oort-cloud comets with crystalline silicates, but these winds could possibly also deposit crystalline silicates in Kuiper-belt comets as well. The low gas densities and orbital velocities of the shock model rule out this possibility for comets that form beyond $\\sim$ 10 AU. If shocks annealed silicate grains {\\it in situ}, the possibility exists that the nebula gas could be chemically processed by the shock as well. For example, Kress et al.\\ (2001) suggest that nebula shocks could produce nitriles that are incorporated into nitrogen-bearing organics in meteorites. Observed correlations of shock-produced, disequilibrium chemicals like nitriles, with silicate crystallinity, would strengthen the case for shock annealing. Comets, nearly all of them long-period comets, despite forming at very cold temperatures, contain crystalline silicate grains that imply temperatures $\\gtsimeq 1100 \\, {\\rm K}$, if only for a fraction of a second. These conditions are met in nebular shocks. Gravitational instabilities make shocks a likely event in the first few Myr of the solar nebula, but whatever the source, silicate grains can be annealed by shocks in the 5 -- 10 AU region of the solar nebula. Shock-annealing obviates the need for large-scale radial transport of annealed grains from the inner solar nebula to the 5 -- 10 AU region where the Oort-cloud comets formed." + }, + "0112/astro-ph0112341_arXiv.txt": { + "abstract": "Using a new radio sample, 6C* designed to find radio galaxies at $z > 4$ along with the complete 3CRR and 6CE sample we extend the radio galaxy $K-z$ relation to $z \\sim 4.5$. The 6C* $K-z$ data significantly improve delineation of the $K-z$ relation for radio galaxies at high redshift ($z >2$). In a spatially flat universe with a cosmological constant ($\\rm \\Omega_{\\rm M}=0.3$ and $\\rm \\Omega_{\\Lambda}=0.7$), the most luminous radio sources appear to be associated with galaxies with a luminosity distribution with a high mean ($\\approx 5L^{\\star}$), and a low dispersion ($\\sigma \\sim 0.5$~mag) which formed their stars at epochs corresponding to $z\\, \\gtsim\\, 2.5$. ", + "introduction": "Radio galaxies provide the most direct method of investigating the host galaxies of quasars if orientation based unified schemes are correct. The nuclear light which dominates the optical/near-infrared emission in quasars is obscured by the dusty torus in radio galaxies, therefore difficult psf modelling and subtraction are not required to determine the properties of the underlying host galaxy. Unfortunately compiling samples of radio loud AGN is a long process, because of the radio selection there is no intrinsic optical magnitude limitation, making follow-up observations extremely time consuming, especially when dealing with the faintest of these objects. However, low-frequency selected radio samples do now exist with the completion of 3CRR (Laing, Riley \\& Longair 1983) along with 6CE (Eales et al. 1997; Rawlings et al. 2001) and the filtered 6C* sample (Blundell et al. 1998; Jarvis et al. 2001a; 2001b). We can now use these radio samples to investigate the underlying stellar populations through the radio galaxy $K-z$ Hubble diagram. ", + "conclusions": "" + }, + "0112/physics0112018_arXiv.txt": { + "abstract": "We show that moderate deviations from the Maxwell-Boltzmann energy distribution can increase deuterium reaction rates enough to contribute to the heating of Jupiter. These deviations are compatible with the violation of extensivity expected from temperature and density conditions inside Jupiter. ", + "introduction": "Jupiter emits more radiation then it recives from the sun: the origin of this excess heat is still uncertain and debated. Possible explanations are: release of gravitational potential energy due to the planet contraction and/or Helium sedimentation \\cite{Guillot:1999bw}, decaying of radioactive isotopes in the core \\cite{Hu:80}, or deuterium burning \\cite{Ho:91}. Each of these hypotheses has difficulties \\cite{Ou:98}; in particular, standard calculations of deuterium burning reaction rates predict negligible contribution to the planet thermal balance, in spite of the substantial enhancement due to electron and ion screening effects \\cite{Ho:91,Ic:93}. In a strongly coupled plasma, anomalous diffusion and time correlation effects originate non-Maxwellian two-body relative energy distribution that can be parameterized with and, in same cased, assume the same functional form that appears in the contest of Tsallis non-extensive thermodynamics \\cite{Kaniadakis:1997,Kaniadakis:1998my}. As demonstrated for the solar core, small changes of the tail of the energy distribution can strongly modify the fusion rates without affecting mechanical properties (hydrostatic equilibrium and the sound speed) that depend on the mean value of the distribution \\cite{Coraddu:1999yb}. Since the internal conditions of Jupiter indicate the existence of a strongly coupled plasma, we investigate the effects of the consequent small deviations from the standard Maxwell-Boltzmann (MB) statistics on deuterium burning rates and the possibility that deuterium burning could play, or have played in the past, a role in Jupiter thermal balance. ", + "conclusions": "In Fig.~\\ref{fig:TsRates} deuterium reaction rates are plotted for two different values of the $q$ parameter. We can observe that moderate deviations $(q-1)\\sim 0.1$ from the Maxwellian distribution increase deuterium burning rates above the threshold rate $r_t \\approx 1 \\textrm{sec}^{-1} \\textrm{cm}^{-3}$: therefore, these processes should contribute to heat Jupiter at the present epoch ($T\\approx 1-2$ eV). If we read the graphs for $T\\approx 10-20$ eV, which corresponds to temperatures of the planet during its formation, we realize that it was sufficient a smaller value, $(q-1)\\sim 0.03$, to effect the thermal balance in that period. Reaction rates of the order of that required to heat Jupiter do not cause decrease significantly the deuterium density inside the planet; in fact a burning rate equal to ten times the threshold rate would consume a significant fraction of deuterium only after a time of the order of $ n_D / (10 r_t) \\sim 10^{11}$ years. These considerations demonstrate that deuterium burning is a possible explanation for the Jupiter excess heat. Precise determinations of the conditions inside Jupiter and additional microscopic calculations are necessary for a better determination of the range of values of $q$ relevant to Jupiter interior." + }, + "0112/astro-ph0112163_arXiv.txt": { + "abstract": "We explore the infrared M band (4.7 \\mum) spectrum of the class I protostar L1489 IRS in the Taurus Molecular Cloud. This is the highest resolution wide coverage spectrum at this wavelength of a low mass protostar observed to date ($R=$25,000; $\\Delta v=$12 \\kms). A large number of narrow absorption lines of gas phase \\twelveco, \\thirteenco, and \\eighteenco\\ are detected, as well as a prominent band of solid \\twelveco. The gas phase \\twelveco\\ lines have red shifted absorption wings (up to 100 \\kms), which likely originate from warm disk material falling toward the central object. Both the isotopes and the extent of the \\twelveco\\ line wings are successfully fitted with a contracting disk model of this evolutionary transitional object \\citep{hoge01}. This shows that the inward motions seen in millimeter wave emission lines continue to within $\\sim$0.1 AU from the star. The amount of high velocity infalling gas is however overestimated by this model, suggesting that only part of the disk is infalling, e.g. a hot surface layer or hot gas in the magnetic field tubes. The colder parts of the disk are traced by the prominent CO ice band. The band profile results from CO in 'polar' ices (CO mixed with H$_2$O), and CO in 'apolar' ices. At the high spectral resolution, the 'apolar' component is, for the first time, resolved into two distinct components, likely due to pure CO and CO mixed with CO$_2$, O$_2$ and/or N$_2$. The ices have probably experienced thermal processing in the upper disk layer traced by our pencil absorption beam: much of the volatile 'apolar' ices has evaporated, the depletion factor of CO onto grains is remarkably low ($\\sim$7\\%), and the CO$_2$ traced in the CO band profile was possibly formed energetically. This study shows that high spectral resolution 4.7 \\mum\\ observations provide important and unique information on the dynamics and structure of protostellar disks and the origin and evolution of ices in these disks. ", + "introduction": "~\\label{sel1489:intro} In the process of low mass star formation, a mixture of gas, dust, and ices accumulates in protostellar envelopes and disks. The fate of this molecular material is diverse. Most of it will fall toward the protostar and dissociates in the inner disk region or stellar photosphere. Some material will be blown away and destroyed by the stellar wind. Some may also survive and be the building material for comets and planets. Major aspects of this complicated process are not well understood, and poorly observationally constrained. For example, do the ices that form comets still resemble ices of the original pristine molecular clouds or are new ices of different composition being formed in the envelope or disk? The type of ices being formed depends on the composition of the gas that accretes onto grains. Reducing environments produce H$_2$O-rich (`polar') ices, while in cold inert environments `apolar' ices rich in CO, N$_2$ and O$_2$ can be formed \\citep{tiel82}. Depending on the composition, ices evaporate between temperatures of 18 and 90 K. Also, heat can change the solid state structure of ices by for example crystallization. Energetic particles (e.g. cosmic rays) and ultraviolet (UV) radiation are able to initiate reactions in ices and form new species. Dynamics and shocks within disks may be able to destroy ices as well. Clearly, to determine the relative importance of these ice formation and destruction processes, knowledge of the physical conditions and structure of envelopes and disks is crucial. Much theoretical and observational work on this topic has been done over the last $\\sim$10 years. Molecular gas was detected in a suite of protostellar disks by millimeter wave observations sensitive to emission over radii of several hundred AU (Dutrey, Guilloteau, \\& Guelin 1997; \\citealt{thi01}). Gas phase abundances were found to be reduced by factors of 5 to several 100, depending on the source and the sublimation temperature of the molecules. Models of disk mid-planes indeed show high depletions because of the formation of icy mantles on grains \\citep{aika97, will98}. The predicted depletions were in fact higher than observed and thus desorption mechanisms are needed to explain the millimeter wave observations \\citep{gold99}. It was realized that the outer parts of disks are heated more efficiently when they are flared \\citep{keny87}. Thus, by the influence of the stellar radiation a layer with `super-heated' dust is formed in which molecules have been dissociated \\citep{chia97}. The layer below that is warm enough to evaporate the ices, but not dissociate the released molecules. The importance of this warm layer, and relative gas phase molecular abundances, depends strongly on how effective ice desorption mechanisms are \\citep{will98}. Recent studies indicate that desorption by UV and X-ray photons may be strong enough to explain observations of molecular gas by millimeter wave telescopes \\citep{will00, naji01}. This idea is confirmed by multi transition molecular line observations which indicate temperatures ($>20-40$ K) and densities that typically occur in this warm layer \\citep{zade01}. In a third layer, the disk mid-plane, cold, dense conditions prevail, resulting in extreme depletions of gas phase molecules many orders of magnitude larger than in quiescent dense molecular clouds. Indeed, recent absorption line observations failed to detect gas phase CO in the edge-on disk around the protostar Elias 18, thus indicating an enormous depletion in the mid-plane \\citep{shup01}. In this Paper we report high spectral resolution ($R=$25,000) 4.7 \\mum\\ M band observations of the obscured protostar L1489 IRS (IRAS 04016+2610) in the Taurus Molecular Cloud. L1489 IRS is a low luminosity object (3.7 $L_{\\odot}$), with a spectral energy distribution resembling that of an embedded class I protostar (Kenyon, Calvet, \\& Hartmann 1987). Detailed millimeter wave line and continuum studies show that L1489 IRS is surrounded by a large, 2000 AU radius rotating thick disk-like structure \\citep{hoge98, sait01}, rather than an inside-out collapsing envelope \\citep{hoge00a}. The rotation is sub-Keplerian, and the disk as a whole is contracting. Thus it was suggested that L1489 IRS represents a short-lasting (2$\\times 10^4$ yr) transitional phase between embedded YSOs that have large envelopes and small (few hundred AU) rotationally supported disks, and T Tauri stars which have no envelopes and fully rotationally supported 500-800 AU size disks \\citep{hoge01}. This circumstellar (or circumbinary: \\citealt{wood01}) disk is seen close to edge-on (60 to $<$90$^{\\rm o}$) in scattered light images \\citep{whit97, padg99}. A CO outflow emanates from the object \\citep{myer88} with Herbig Haro objects lying along it \\citep{gome97}. Low spectral resolution infrared observations of L1489 IRS show that deep H$_2$O \\citep{sato90} and CO (\\citealt{chia98}; Teixeira, Emerson, \\& Palumbo 1998) ice bands are present along the line of sight. In this Paper we will use the newly available spectrometer NIRSPEC at the Keck II telescope to obtain high resolution M band spectra ($\\Delta v=$12 \\kms\\ at 4.7 \\mum) of this source. The large array of NIRSPEC allows both the vibrational band of solid CO and the surrounding ro-vibrational transitions of gas phase CO to be observed in the same high resolution spectrum. This offers a new view on this system, both on the origin and evolution of ices and the interrelationship of gas and ices as well as on the kinematics and structure of the young, contracting, close to edge-on disk. It is a unique view, because infrared absorption line studies trace all gas and solid state material at all radii from the star, while present day millimeter wave observations are limited by their relatively low spatial resolution ($\\geq$100 AU). Thus one of the questions that will be answered in this Paper is whether the large scale inward motions seen in millimeter wave emission lines continue to smaller radii (1 AU or less) from the star. Previous studies have already shown that rich astrophysical information can be obtained from high spectral resolution observations in the atmospheric M band. Mainly massive, luminous protostars were observed, however \\citep{mitc90}, and the few observations of low mass protostars cover a small wavelength range containing only a few gas phase lines and not the solid CO band \\citep{shup01, carr01}. This is Paper I in a series on high resolution M band spectroscopy of protostars, initiated by the availability of the NIRSPEC spectrometer at Keck II with which weak, low mass protostars can be routinely observed at high spectral resolution over a large wavelength range covering both the solid and gas phase CO features. The reduction of the long slit spectra is discussed in \\S 2. In \\S 3.1, we analyze the observed gas lines, which at this resolution even give dynamical information. We use standard curve of growth and rotation diagram techniques to get a first idea of gas column and temperatures. In order to analyze the solid CO band profile in this line of sight, a detailed discussion of available laboratory experiments of solid CO is given in \\S 3.2. In \\S 4.1 we apply the infalling disk model of \\citet{hoge01} to explain the observed gas phase \\twelveco\\ and \\thirteenco\\ line profiles, and constrain the physical conditions and structure of the disk. The possibility of binarity is briefly discussed in \\S 4.2. The gas phase analysis is linked to the the solid CO results to determine the origin and thermal history of solid CO in \\S 4.3. We conclude with suggestions for future work in \\S 5. ", + "conclusions": "\\subsection{An Infalling Disk} The astrophysical meaning of the apparent two component temperature structure seen in the \\thirteenco\\ rotation diagram (Fig.~\\ref{f:rot}) requires further investigation. For high mass protostars it was found that similar rotation diagrams can be `mimicked' by power law models of spherical envelopes \\citep{tak00}. For L1489 IRS, the detection of molecular gas at a range of temperatures and red shifted velocities could indicate the presence of infalling gas at a range of radii from the protostar. Indeed, a 2000~AU radius contracting, disk-like structure was found in millimeter wave interferometer data \\citep{hoge00a}. In a detailed follow-up study, \\citet{hoge01} adopts a flared-disk model based on \\citet{chia97} with a radial power-law distribution for the temperature \\begin{equation} T = 34(R/1000\\,{\\rm AU})^{-0.4} {\\rm ~K,} \\end{equation} and a density distribution that has a power-law drop-off with radius and a vertical exponential drop-off with scale height $h$ \\begin{equation} \\rho(R,z) =\\rho_0 (R/1000 {\\rm AU})^{-1.5} {\\rm e}^{-z^2/h^2} {\\rm ~kg~cm^{-3}}. \\label{eq:den} \\end{equation} The scale-height $h$ is assumed to be a simple function of R, $h=R/2$. An inward-directed radial velocity field described as \\begin{equation} V_{\\rm in} = 1.3 (R/100\\,{\\rm AU})^{-0.5} {\\rm ~km~s^{-1}} \\end{equation} is inferred, in addition to Keplerian rotation around a 0.65~M$_\\odot$ central star. \\begin{figure*}[t!] \\center \\includegraphics[angle=90, scale=0.72]{f6.eps} \\caption{Comparison of the observed spectrum of L1489 IRS with an infalling disk model. The histograms in panels (a) and (b) represent the average of the observed \\twelveco\\ P(6)--P(15) lines and all \\thirteenco\\ lines respectively, corrected for the source and earth velocity (43 \\kms). The smooth thick gray line in these two panels is the collapsing disk model, averaged over the same lines. In panels (c) and (d) the same averaged observed spectra and model are plotted, but now assuming that in the model 1\\% of the unextincted stellar continuum emission does not pass through the disk. Panel (e) shows the latter model (thick gray line) compared to a portion of the non-averaged L1489 IRS spectrum.}~\\label{f:avemod} \\end{figure*} Can this contracting disk model, based on (sub-) millimeter emission observations with angular resolution of $4$--$8''$, reproduce the observed infrared CO absorption line profiles measured along a pencil beam? The absorption lines are modeled with the radiative transfer code of \\citet{hoge00b}; the high densities in the disk ensure LTE excitation for the lines involved, and line trapping is neglected in the excitation calculation. The model spectra include dust opacity at a standard gas/dust ratio, as well as a $N({\\rm CO})=1\\times 10^{18}$ cm$^{-2}$ column of cold foreground material (15 K; \\citealt{hoge01}); both factors do not affect the spectra in any significant way. The calculated spectrum is convolved with a Gaussian of FWHM=12 \\kms, which is the NIRSPEC instrumental resolution. We find that, while keeping all other parameters the same as in \\citet{hoge01}, the assumed density profile sensitively influences the wings of the $^{12}$CO lines. This is enhanced by the fact that we are observing the flared disk of L1489 IRS at an inclination between $60^\\circ$ and $<90^\\circ$ (cf., \\citealt{padg99}), and the pencil beam crosses the disk at a few scale heights. Small changes in the density profile, for example induced by the thermal structure, have a large effect on the absorption line profile. In the model of \\citet{hoge01} the scale height increases linearly with distance from the star, and thus the density $\\rho (l)$ along the line of sight $l$ follows the density in the mid-plane (Eq.~\\ref{eq:den}) reduced by a factor $e^{-4/\\tan^2(\\alpha)}$, with $\\alpha$ the inclination. Here, we include the effect of density variations, or deviations from the adopted scale height $h=R/2$, by relaxing the values of the density along the line of sight $l$ by fitting $\\rho(l) = \\rho_0 (l/1000 {\\rm ~AU})^{-p}$ to the data. This initial model successfully fits the peak velocity and depth of both high and low $J$ $^{13}$CO lines (Fig.~\\ref{f:avemod}). Its rotation diagram is quite different from that of the curve of growth analysis (Fig.~\\ref{f:rot}), showing that rotation diagrams must be interpreted with great care. Our model also matches the range of velocities observed in the red wings of the $^{12}$CO lines, when taking $p=0.55\\pm 0.15$. This is a much shallower density profile compared to that derived from millimeter wave data ($p=1.5$; Eq.~\\ref{eq:den}) and indicates that the scale height increases more than linear, i.e. the disk flares more than assumed in \\citet{hoge01}. With this result, it is possible to determine the important relation of disk scale height $h(R)=a.R^b$ as a function of $R$, but only if the disk inclination is {\\it a priori} known. Unfortunately the inclination is not better constrained than within the range of $60^\\circ$ and $<90^\\circ$ imposed by near-infrared data (\\citealt{padg99}). We can therefore not distinguish between low and high values of $a$ and corresponding high and low inclinations respectively. In either case, the total $^{12}$CO column along the pencil beam is $1.2\\times 10^{19}$ \\sqcm, with 58\\% of the CO mass at a temperature of $T=20-60$ K, 15\\% at $60-90$ K, and 27\\% at $60-90$ K. This result is of importance in \\S 4.3 in the interpretation of the solid CO observations, and in particular in assessing the thermal history of ices. The total column of our model is in good agreement with the column derived from the visual extinction ($1.4\\times 10^{19}$ \\sqcm; \\S 3.1). It is also of the same order of magnitude as the total column through the mid-plane, calculated from dust and line emission ($N[{\\rm CO}]=6\\times 10^{18}$ \\sqcm; \\citealt{hoge01}), and confirms the relatively edge-on orientation of the disk. However, apart from these successes, the \\twelveco\\ lines show that our infalling disk model produces too much warm gas at high velocities (Fig.~\\ref{f:avemod}). The $^{12}$CO lines are a factor of 2.5 deeper, and, in contrast to $^{13}$CO, they peak at a too high velocity (+10 \\kms) with respect to the observations. In principle, one could make the \\twelveco\\ lines less deep by assuming that $\\sim 1\\%$ of the original, unextincted continuum flux (corresponding to 30\\% of the extincted continuum) reaches the slit without passing through the disk, by scattering on large grains. The shift in peak velocity however requires a solution of a more fundamental origin. Perhaps the infall velocity function is shallower, and the disk is more rotationally supported at lower radii. The amount of warm gas at high velocities can also be lowered by assuming that only part of the disk participates in the high velocity inflow, such as a thin hot surface layer, or gas accelerated in magnetic field tubes directed from the inner disk to the stellar photosphere. Such a two component model is consistent with the rotation diagram derived from the curve of growth (Fig.~\\ref{f:rot}), and also with the low observed mass accretion rate. If we take the inflow at face value, and assume that the entire disk participates, the mass accretion rate would be $10^{-6}$ M$_\\odot$, generating 7 L$_\\odot$ in accretion luminosity. The star's L$_{\\rm bol}$ is estimated at 3.7 L$_\\odot$ which also contains the stellar luminosity. It is therefore indeed likely that the mass accretion onto the star is significantly lower, as is also traced through the lack of the hydrogen Pf$\\beta$ emission line in our spectrum (2148.8 \\waven; Fig.~\\ref{f:obs}) and the weakness of Br$\\gamma$ emission \\citep{muze98}. \\subsection{Binarity?} An entirely different explanation for the line profiles may lie in the possibility that L1489 IRS is a protobinary system. A protobinary nature of L1489 IRS is suggested by various pieces of evidence (\\citealt{luca00}; \\citealt{wood01} and references therein). The presence of a quadrupolar outflow system is inferred from K band polarization images, C$^{18}$O emission line profiles, Herbig-Haro knots that are scattered throughout the L1489 IRS environment, and a very complex near infrared scattered light pattern. Three dimensional models, in which the axisymmetry of the infalling circumstellar envelope is broken by multiple outflow cavities that are perpendicular to each other, are able to account for the observed morphology. The putative binary itself, however, has not been resolved so far. An upper limit on the projected separation from near infrared images has been set on $<$ 20 AU \\citep{padg99}. If the CO absorption line profile is in any way related to a binary system, then the large observed velocities ($\\sim$23 \\kms; \\S 3.1) may indeed favor a close binary system. The line profile is then expected to vary on a time scale of a few months, which can easily be tested. In this case, much of the observed warm gas might be present in two small circumstellar disks, which are in a close orbit around each other. Some of the warm gas may also be present at low density in the central cavity created by the binary. The large column of cold gas may originate in the circumbinary disk. The binary extracts momentum from the 2000 AU circumbinary disk, setting up the inward motion seen in millimeter wave emission lines. We leave further investigation of this topic for future studies. \\subsection{The Origin and Evolution of Ices} In order to establish if the solid CO observed toward L1489 IRS originates in foreground clouds or in a circumstellar (or binary) disk, it is worth to compare with ices observed in lines of sight not affected by star formation. Observations of field stars obscured by intervening quiescent material of the Taurus Molecular Cloud have revealed that solid CO is not present when the extinction $A_{\\rm V}\\lesssim 5$ (e.g. \\citealt{teix99}). The solid CO toward L1489 IRS can therefore not be associated with foreground clouds, which have a gas column of $N$($^{12}$CO)= 1$\\times 10^{18}$ \\sqcm\\ \\citep{hoge01}, corresponding to $A_{\\rm V}\\sim 2$. Thus, the solid CO must be present in the disk of L1489 IRS. The absorption profile is intriguingly different from that seen in quiescent clouds. The broad red wing has a depth of $\\sim$30\\% with respect to the narrow 2140 \\waven\\ peak, which is significantly more than toward all measured background stars (10\\%; \\citealt{chia95}). This may well be an effect of thermal processing along the L1489 IRS line of sight, because the sublimation temperatures of polar and apolar ices, causing the broad and narrow features respectively, are very different (90 versus 18 K). However, a chemical origin of an increased abundance of polar ices in disks cannot be excluded, because the apparently edge-on system Elias 18 in the Taurus Molecular Cloud has an extremely large CO depletion factor (solid/[gas+solid]$\\sim$100\\% versus 7\\% for L1489 IRS), but a deep red `polar' CO wing is present as well \\citep{shup01, chia98}. On the other hand, energetic processing may take place even in the cold disk of Elias 18 \\citep{whit01}. Clearly, it is necessary to observationally characterize the ices in circumstellar disks in much more detail. If for now we assume the sublimation scenario, we can do some general extrapolations which can be compared with the results of our gas phase study (\\S 4.1). By scaling the long wavelength wing of solid CO of background field stars to that of L1489 IRS, we find that a column of 6$\\times 10^{17}$ \\sqcm\\ of CO has evaporated from the apolar ice component in the part of the L1489 IRS disk along the pencil absorption beam where $T<90$ K (the sublimation temperature of polar ices). Then the column of solid CO that went from the quiescent cloud into building this part of the disk is 12.5$\\times 10^{17}$ \\sqcm. Extrapolating this further, we use the observed CO depletion factor of 30\\% toward field stars behind the Taurus Molecular Cloud \\citep{chia95} to calculate that the original quiescent gas column must have been of the order of 3$\\times 10^{18}$ \\sqcm. Adding the evaporated column, the expected present day gas column at $T< 90$ K is 3.6$\\times 10^{18}$ \\sqcm. This is of the same order of magnitude as the CO column below 90 K in our collapsing disk model ($N$[CO]=8.7$\\times 10^{18}$ \\sqcm), which may indicate that no chemical change in the apolar/polar CO ice ratio and no significant additional depletion has occurred in the evolution from quiescent Taurus Molecular Cloud material to the formation of the L1489 IRS disk. This contrasts strongly with the very large depletions found in the (older) disks of T Tauri stars \\citep{dutr97}. The low CO depletion along the pencil beam toward L1489 IRS (7\\%) and the supposed signs of thermal processing (see below) may be due to the fact that our line of sight does not cross the disk mid-plane, i.e. the system is not exactly edge-on. The ice processing we see takes place higher in the disk atmosphere, perhaps in the warm layer below the super-heated dust layer responsible for millimeter wave line emission \\citep{zade01}. It must be noted that in the model of \\citet{hoge01} the gas temperatures are larger than 25 K, prohibiting the formation of apolar ices and large CO depletions anywhere in the disk. The observed presence of apolar CO ices thus indicates that, as already suggested in \\S 4.1, the line of sight may cross the cold, rotationally supported disk interior not traced in the observations and infall model of \\citet{hoge01}. Apart from evaporation of apolar ices, other hints of thermal processing include the aforementioned absence of the 2150 \\waven\\ absorption (\\S 3.2), which occurs in cold unprocessed polar CO ices but disappears at temperatures $T>50$ K. Also, the blue apolar wing may be a consequence of thermal processing. If the central 2140 \\waven\\ peak is due to a mixture of O$_2$, N$_2$ and CO instead of pure CO (spectroscopically these cannot be distinguished), thermal or energetic processing (UV radiation from the ISRF, UV induced by H$_2$ cosmic ray collisions, or direct hits of cosmic rays) could efficiently produce CO$_2$. This could cause the band to broaden and shift to the position of the observed blue wing. Chemical models indicate that energetic processing of molecules in disks takes place on a time scale of 10$^6$ yrs \\citep{aika99}, which is somewhat longer than the age of the disk of L1489 IRS ($\\sim 5\\times 10^5$ yrs). This however applies to the disk mid-plane, and the time scale may well be shorter in the lower density higher disk layers that our observations of L1489 IRS trace. A possible problem with the energetic processing interpretation is the absence of a feature adjacent to the short wavelength side of the CO ice band, usually attributed to energetically produced C$\\equiv$N bondings \\citep{whit01}. Another spectroscopic tracer of thermal processing is the signature of crystallization in the band profiles of H$_2$O and CO$_2$ ices. Our infalling disk model predicts that only 15\\% of the gas is within the temperature range at which ices crystallize (60--90 K), and thus crystallization is not expected to play a significant role in the disk of L1489 IRS. This model prediction can be tested with future high quality H$_2$O and CO$_2$ spectra of L1489 IRS. In summary, several pieces of evidence indicate that the CO ices in the disk of L1489 IRS have experienced thermal or energetic processing. The strongest arguments are the low depletion factor and the low ratio of apolar to polar ices with respect to the quiescent Taurus Molecular Cloud material. This may be explained by the fact that the disk of L1489 IRS is seen under an angle, and our pencil absorption beam traces the warm upper disk layers. \\vspace{15pt}" + }, + "0112/astro-ph0112449_arXiv.txt": { + "abstract": "It has become increasingly apparent that traditional hydrodynamical simulations of galaxy clusters are unable to reproduce the observed properties of galaxy clusters, in particular overpredicting the mass corresponding to a given cluster temperature. Such overestimation may lead to systematic errors in results using galaxy clusters as cosmological probes, such as constraints on the density perturbation normalization $\\sigma_8$. In this paper we demonstrate that inclusion of additional gas physics, namely radiative cooling and a possible preheating of gas prior to cluster formation, is able to bring the temperature--mass relation in the innermost parts of clusters into good agreement with recent determinations by Allen, Schmidt \\& Fabian using {\\it Chandra} data. ", + "introduction": "Reproducing the observed number density of rich galaxy clusters has long been thought to be one of the most reliable constraints on the matter power spectrum on short scales. It has been studied by many authors over the years (\\citealt{E89}; \\citealt{HA91}; \\citealt*{WEF93}; \\citealt*{EKF96}; \\citealt{VL96}, \\nocite{VL99}1999; \\citealt{H97}, \\nocite{H00}2000; \\citealt{BSB00}; \\citealt*{PSW01}; \\citealt{W01}), recent determinations typically yielding $\\sigma_8 \\sim 0.9$ to $1.0$ for the currently-favoured $\\Lambda$CDM model with matter density $\\Omega_0 \\simeq 0.3$. However, recently evidence has begun to accumulate from a number of sources that this may be a significant overestimate, perhaps by tens of percent. For example, the required $\\sigma_8$ estimated from the 2dF galaxy survey (\\citealt{L02}; \\citealt{V02}), or from that survey combined with other probes (\\citealt{E02}), is significantly lower, and there are now several papers using galaxy clusters that also give lower results (\\citealt{BRT01}; \\citealt{RB02}; \\citealt*{VNL02}). A low value was also found recently by \\citet{S02}, who used an observed relation between cluster temperature and mass \\citep*{FRB01} rather than one derived from hydrodynamical simulations. This last result is particularly significant, and points to the increasingly evident result that traditional hydrodynamical simulations, which include only adiabatic gas heating during collapse, are unable to reproduce the observed properties of clusters. For example, the recent {\\it Chandra} results of \\citet*[][hereafter ASF01]{ASF01} indicate that, at least in the inner regions where data exists, clusters are considerably hotter for a given mass than predicted by adiabatic simulations. Here we address the question of whether the inclusion of additional gas physics, both radiative cooling of the gas and preheating of the gas before cluster formation, is capable of bringing the simulations into agreement with observations. We concentrate only on the inner regions of clusters, for which temperature profiles have been measured by {\\it Chandra}, and we find that indeed the observations can be reproduced. In itself this is not sufficient aid to theorists seeking to constrain $\\sigma_8$, which requires an accurate description of clusters out to the virial radius, but this encouraging result suggests that simulations may soon be useful for this purpose. We will explore the cluster temperature--mass relation out to larger radii in a forthcoming paper. ", + "conclusions": "We have shown that simulations are capable of reproducing the observed relationship between mass and temperature in the inner regions of galaxy clusters. In particular, the mass-weighted temperature versus mass within a radius enclosing an overdensity of 2500 in our Radiative and Preheating simulations agrees with the observed relation of ASF01. There are a number of caveats, however. The temperature range of the simulations and the observations barely overlap; we have one cluster above 6\\,keV, while ASF01 have only one below this temperature. Nevertheless, there is no reason to suppose that our results will not extend up to higher temperatures, though confirmation of this will have to await resimulations of clusters drawn from larger simulation boxes. Perhaps more pertinently, none of our simulations presents a fully realistic model of clusters, the Radiative model producing too much cooled gas and the Preheating model too little. However they both match the observed X-ray luminosity--temperature relation, because they both have a higher entropy within \\rdel\\ than does the Non-radiative simulation. This increase in entropy manifests itself as an increase in the temperature of the gas in the inner parts of the clusters. One might expect, therefore, that realistic clusters that share the same entropy profile would predict the same temperature--mass relation. Unfortunately, the results presented in this paper and in ASF01 are of limited use to theorists who wish to predict the temperature function of clusters in order to constrain cosmology. This is because they need to relate the mass within the virial radius to the emission-weighted temperature of clusters. The prediction of masses at $r_{500}$ or larger radii from the X-ray observations is a harder problem than discussed here and will be investigated in a longer paper." + }, + "0112/astro-ph0112213_arXiv.txt": { + "abstract": "In this paper we present and analyze new CCD $UBVRI$ photometry down to $V~\\approx$~21 in the region of the young open cluster Trumpler~15, located in the Carina spiral feature. The cluster is rather compact and has a core radius of about 2$^{\\prime}$, which translates in about 1 pc at the distance of the cluster. We provide the first CCD investigation and update its fundamental parameters. We identify 90 candidate photometric members on the base of the position in the color-color and color-magnitude diagrams. This sample allows us to obtain a distance of 2.4$\\pm$0.3 kpc from the Sun and a reddening E$(B-V)$~=0.52$\\pm0.07$. We confirm that the cluster is young, and fix a upper limit of 6 million yrs to its age .\\\\ In addition, we draw the attention on the lower part of the Main Sequence (MS) suggesting that some stars can be in contracting phase and on a gap in the MS, that we show to be a real feature, the $B1-B5$ gap found in other young open clusters.\\\\ We finally study in details the extinction toward Trumpler~15 concluding that it is normal and suggesting a value of 2.89$\\pm$0.19 for the ratio of total to selective absorption $R_V$.\\\\ ", + "introduction": "In this paper we study the stellar content of the young compact open cluster Trumpler~15 by means of deep multicolor CCD photometry.\\\\ This open cluster ($\\alpha$~=~10:44:33.0, $\\delta$~=~ -59:24:24.0, $l$~=~287.41, $b$~=~ -0.41; J2000.0) is located near the northern edge of the Great Carina Nebula (NGC~3372), about $20^{\\prime}$ above $\\eta$~Carin\\ae~. It is also named VdB-Hagen~104, Lund~558 and OCL~825.\\\\ Like other young clusters in this region (e.g. Trumpler~14 and 16), is rather compact and rich. There are several intriguing questions related to this cluster, which was discussed in the past often leading to contradictory results. Is the interstellar extinction toward Trumpler~15 normal? Is this cluster connected with the other ones located much closer to $\\eta$~Carin\\ae~, like Trumpler~14, 16 and Collinder~232? In other words, does it share the same properties of these clusters, like age and distance, suggesting that it probably formed together with them in the same Star Formation event?\\\\ \\noindent Aiming at clarifying these issues and deriving updated estimates for its fundamental parameters, like distance and age, in this paper we present and discuss the first $UBVRI$ CCD photometric study of Trumpler~15.\\\\ The layout of the paper is as follows: Section~2 presents briefly the data acquisition and reduction. In Section~3 we discuss previous investigations on this cluster; in Section~4 we compare our photometry with previous ones and present our data. Section 5 illustrates the technique to derive reddening and membership of stars in Trumpler~15. Section~6 is dedicated to the study of the interstellar extinction toward the cluster, while in Section~7 we derive estimates for Trumpler~15 age and distance. Finally, Section~8 discusses the geometrical structure of the cluster and Section~9 summarizes our findings. \\begin{figure*} \\centerline{\\psfig{file=fig1.eps,width=17cm,height=17cm}} \\caption{A map of a observed region around Trumpler~15. The size of each star is proportional to its magnitude. North is up, East on the left. The field is about $6^{\\prime} \\times 6^{\\prime}$.} \\end{figure*} ", + "conclusions": "In this paper we have presented new $UBVRI$ CCD photometry for Trumpler~15, a young open cluster located in the Carina spiral feature.\\\\ We identify 90 photometric members on the base of individual reddenings, position on the CMDs and spatial distribution in the field.\\\\ Basing on this large sample we provide updated estimates of cluster fundamental parameters. We find that the cluster is young, with an age between 2 and 6 million yrs, contains a possible population of pre MS candidates, which deserves further investigation, and shows a gap in the MS at $V_o \\approx 10.5$, that we suggest to be a real feature (the $B1-B5$ gap already found in other clusters).\\\\ We place the cluster at $2.4\\pm0.3$ kpc from the Sun. Moreover we obtain E$(B-V)$~=~0.52$\\pm$0.07, and find that the extinction toward Trumpler~15 can be considered normal. Finally we estimate that the cluster has a core radius of about $2^{\\prime}$.\\\\ Trumpler~15 appears to be located somewhat closer to the Sun than Trumpler~14, Trumpler~16 and Collinder~232. Nevertheless, the data suggest that Trumpler~15 might belong to the complex defined by Trumpler~14, Trumpler~16, Collinder~232 and Collinder~228, since it shares with these clusters the same age and the presence of pre MS candidates. Moreover the extinction law seems to be basically normal (with some local fluctuations) in the entire region.\\\\ \\noindent The most appealing scenario one can envisage is that all these clusters probably formed together in the same recent Star Formation event." + }, + "0112/astro-ph0112025_arXiv.txt": { + "abstract": "{We present a new method to estimate the Hubble constant $H_0$ from the measured time delays in quadruply imaged gravitational lens systems. We show how it is possible to get an estimate of $H_0$ without the need to completely reconstruct the lensing potential thus avoiding any {\\it a priori} hypothesis on the expression of the galaxy lens model. Our method only needs to assume that the lens potential may be expressed as $r^{\\alpha} F(\\theta)$, whatever the shape function $F(\\theta)$ is, and it is thus able to fully explore the degeneracy in the mass models taking also into account the presence of an external shear. We test the method on simulated cases and show that it does work well in recovering the correct value of the slope $\\alpha$ of the radial profile and of the Hubble constant $H_0$. Then, we apply the same method to the real quadruple lenses PG1115+080 and B1422+231 obtaining $H_0 = 58^{+17}_{-15} \\ {\\rm km \\ s^{-1} \\ Mpc^{-1}}$\\,(68 \\% CL). ", + "introduction": "Most of the ways of measuring the Hubble constant $H_0$ involve a form of distance ladder, which utilizes a number of astrophysical standard candles and standard ruler relations, and are calibrated locally by a geometrical technique such as trigonometric or dynamical parallaxes (e.g., \\cite{M98}). A few methods involve no distance ladder\\,: good examples are (i) inferring the distance of the SNeII from their light curves and spectra by modelling their expanded photosphere (\\cite{S92}) and (ii) comparing the $H_0$\\,-\\,independent angular extent of galaxy clusters to their $H_0$\\,-\\,dependent depth as deduced by the X\\,-\\,ray emission and the Sunyaev\\,-\\,Zeldovich microwave background decrement due to the cluster itself (\\cite{HB98}). But the most promising one step method may be considered the one first proposed by Refsdal in 1964 and became feasible only recently thanks to the now available instrumentation. The principle of the Refsdal method is quite simple. If a QSO is multiply imaged by the gravitational lensing effect of a galaxy along the line of sight, the light rays coming from the different images follow different optical path and thus arrive to the observer with a time delay among each other. It is easy to show that these time delays are proportional to the inverse of the Hubble constant $H_0$ and to a factor which depends only on the lensing potential and the source coordinates. Having measured the time delays and estimated the lens dependent factor by the images configuration, one can then obtain a direct estimate of the Hubble constant avoiding all the problems and possible systematic errors connected to the distance ladders. There are actually more than fifty multiply imaged systems (both double and quadruple) and the number of them with measured time delays is increasing (\\cite{S00}) so that the prospects of obtaining an accurate estimate of $H_0$ from gravitational lenses is quite good (\\cite{K01}). However, there is still a major problem connected to the modelling of the lensing galaxies since there are often different models which predict the same images configuration and the other lensing observables. Thus lens modelling is the major source of uncertainty in the Refsdal method. There are two ways to compensate for our lack of knowledge about the lens galaxy. The first is assuming an exact parametric form for the galaxy model and then determine its parameters by fitting to the images positions and time delays ratios (and, eventually, to the flux ratios). However, it is clear that this approach strongly underestimates the uncertainty connected to the lens modelling. Even if one is able to find a parametric galaxy model which is dynamically possible and which reproduces the image properties with acceptably low $\\chi^2$, one still has to aggressively explore all other classes of models to get the true uncertainty on $H_0$. But this is not possible with the parametric approach since there is an unavoidable limit to the number of parameters which can be used to describe the galaxy fixed by the number of available constraints. Thus parametric techniques may explore only the simplest models, i.e. one is restricted to a narrow area in the space of the models. To explore it in a systematic fashion one has to follow the second approach using a representation of the galaxy which is as general as possible and thus not restricted to a particular form. One possibility in this sense is to pixelate the galaxy map and consider each pixel as an independent mass element. This is what is done in the {\\it pixellated lens method} (\\cite{SW97}; \\cite{WS00}) which indeed is a valide alternative to the usual parametric techniques. Introducing as less as possible constraints on the reconstructed mass distributions, the pixellated lens method is very efficient in exploring the models space, but it has also the risk of overestimating the uncertainties on $H_0$ since one has almost no way to control if the reconstructed models are physically meaningfull or not. To overcome this difficulty we have elaborated a new approach which aims at estimating the Hubble constant from a detailed exploration of that region of the models space which is compatible with the lensing observables and, at the same time, is physically well motivated. To do this we have to lose some generality since we introduce a lensing potential which has a well defined radial profile, but we still do no hypotheses on the angular part and take also into account the presence of the external shear. Even if less general than the pixellated lens method, our semianalytical technique may be considered as a compromise between the usual parametric technique and the full non\\,-\\,parametric approach. We do not introduce any defined galaxy lens model, but we take into account all the models which give rise to a lensing potential of the form $r^{\\alpha} F(\\theta)$, whatever the shape function $F(\\theta)$ is, and finally select only the ones which reproduces the lensing observables and are physically well motivated. Thus our method is still able to carefully explore the models space and, at the same time, it does not introduce any overestimate connected to the inclusion of unphysical models. The plan of the paper is as follows. In Sect. 2 we do some general considerations on the lens equations and write down the relations which will be used to build the method. This is presented in Sect. 3 where we will show how it is possible to algebrically manage the lens equations to finally get a system which is numerically solvable. Since the set of equations is nonlinear, we have elaborated a simple algorithm that allows us to recover all the solutions and exclude the unphysical ones. The code and the selection criteria are then presented in Sect. 4, while the following Sect. 5 is devoted to the application of the method to simulated cases and to the description of how we extract the Hubble constant estimate from the set of solutions. Having so checked that the method indeed works, we can now apply it to real systems; this is the subject of Sect. 6 where we discuss the real lenses PG1115+080 and B1422+231 and obtain our final estimate of $H_0$. Sect. 7 is then devoted to conclusions. ", + "conclusions": "In this paper we have presented a new semianalytical method to estimate the Hubble constant from the measured time delays in quadruply imaged gravitational lens system. Assuming that the galaxy lens potential may be splitted into a radial part described by a simple power\\,-\\,law profile and an angular part described by a quite general shape function, we have been able to write down a nonlinear set of equations taking into account also the contribution of the external shear. This system may be solved numerically allowing us to fully take into account the lens models degeneracy; a set of physical constraints is then used to select the only reasonable solutions thus avoiding the risk of including in our considerations also non physically motivated models. The final set of solutions may be seen as a parametrization of that class of lens models which are able to generate the observed images configuration and the measured time delay ratio. The class of models so delineated may be translated in a sample of values for the Hubble constant $H_0$ which leads us to the final estimate of this cosmological parameter. To test the method we have applied it to three different simulated lens systems varying both the angular and the radial part of the lens potential and also the external shear parameters. This tests have shown us that the method indeed works and it is also very efficient in determining the slope $\\alpha$ of the radial profile with a reasonable accuracy. Even if this were not our final aim, the ability of our method to find out the $\\alpha$ parameter is a very interesting byproduct since this quantity is very usefull in modelling the dark halos which may be considered as the most important galaxy component responsable of the observed quadruply imaged lens system. The tests have also shown us that the method is not efficient in recovering the Hubble constant; even if the central value of the 90 \\% range individuated for $H_0$ is very near to the input value, the range itself is quite large. This is not an unexpected result since we are fully taking into account the lens models degeneracy and it is well known that this increases the uncertainty on the Hubble constant. Motivated by this result, we have tried to reduce the uncertainty on $H_0$ combining the estimates from different systems. To this aim we simply build a combined histograms of the $H_0$ values (binned by $0.1 \\ {\\rm km \\ s^{-1} \\ Mpc^{-1}}$) multiplying the probabilities from each system and then excluding from the final sample those values of $H_0$ which turn out to have a final probability less than 1\\%. This procedure allows us to reduce the uncertainty on $H_0$ by combining the three simulated systems. The final estimate of $H_0$ is perfectly consistent with the input value used for the simulations which confirms us that the method indeed works. Given these encouraging results, we have then applied our method to the real lenses PG1115+080 and B1422+231, which are two of the only three quadruple lenses for which the time delay has been measured. As regard PG1115+080, we find a slope $\\alpha = 1.03_{-0.20}^{+0.24}$ (90\\% CL) which indicates a near isothermal model consistent with previous models in literature. The 90\\% range for $H_0$ turns out to be quite large ($25 \\div 78 \\ {\\rm km \\ s^{-1} \\ Mpc^{-1}}$), but we note that the central value ($H_0 = 68 \\ {\\rm km \\ s^{-1} \\ Mpc^{-1}}$) is consistent with the values quoted in literature and obtained with different method. As regard B1422+231, we find $\\alpha = 1.08_{-0.36}^{+0.33}$ and $H_0 = 25 \\div 78 {\\rm km \\ s^{-1} \\ Mpc^{-1}}$, with a distribution of values for the Hubble constant peaked towards lower ones. As we know, this is the first time that this system has been taken in consideration to determine the Hubble constant since the time delays have been measured only recently. Combining the resulting histograms from PG1115+080 and B1422+231 leads us to the following final estimate for the Hubble constant\\,: \\begin{displaymath} H_0 = 58_{-15}^{+17} \\ {\\rm km \\ s^{-1} \\ Mpc^{-1}} \\ (68\\% CL) \\ , \\end{displaymath} or more conservatively : \\begin{displaymath} H_0 = 56_{-26}^{+22} \\ {\\rm km \\ s^{-1} \\ Mpc^{-1}} \\ (90\\% CL) \\ . \\end{displaymath} We note that the uncertainty singnificantly increases when passing from the 68\\% to the 90\\% range for the estimate of $H_0$ and that the magnitude of the increasement is higher than expected when compared with the result from our simulations. Analyzing in detail the data, however, it is easy to see that this strange behaviour is completely due to the histogram obtained for B1422+231, which predicts too many models with low values of $H_0$. A possible explanation of this strange behaviour may be connected to the very high uncertainty on the measured time delays which translates to an high uncertainty on the recovered parameters $\\alpha$ and $h$. Should we have chosen different values for the time delays (still within the uncertainties), we should have obtained lower values for $\\alpha$ and higher values for $H_0$ thus narrowing the 90\\% range. Some tests have suggested us that this could be actually a possible explanation. However, also considering only the 68\\% CL, our final estimate of $H_0$ is consistent with all the previous estimates in literature whatever is the method used and the estimators considered. {The new semianalytical method presented here turns out to be a valid complementary alternative to the usual parametric approach and to the fully non\\,-\\,parametric techniques. Parametric methods are usefull in building up galaxy models which may be easily compared to other galaxy observables (if possible) also not directly connected with the lensing charachteristics. However, they have only a modest power in exploring the parameter space since one cannot include too many parameters in the model to avoid having a number of degrees of freedom too low in the $\\chi^2$ minimization. Given that the number of constraints is fixed (eight from images positions and two from time delays ratios if the system is a quadruple one), there in an unavoidable limit to the accuracy of the galaxy model and to the possibility to explore the wide range of lens models that may fit the same lensing observables. On the other hand, non parametric methods try to introduce as less {\\it a priori} hypotheses as possible thus fully exploring the space of the models. However, a fully non parametric approach as the one adopted in the pixellated lens method of Williams \\& Saha (2000) does not allow to control if the models considered are physically motivated or not thus risking to overestimate the uncertainties connected to the modelling problem. Our new method is less general than the pixellated one, but it has the advantage that it selects only models which are physically reliable. Besides, combining the results from different quadruple lens systems helps in reducing the other unidentified sources of systematic errors connected to single systems leading to a final estimate for $H_0$ which correctly takes into account all the possible sources of errors. Further improvements are however still possible. Numerical simulations of galaxy formation in different cosmological backgrounds have suggested that the dark halo density profile may be described by a simple universal law (see, e.g., \\cite{NFW97}, \\cite{Moore}, \\cite{Klypin}). These models lead to lensing potentials which may not be described by Eq.(\\ref{eq: psilens}) since the radial profile is not a single power\\,-\\,law, having different radial slopes for the inner and the outer parts. This could suggest that the potential we have used is not realistic and thus our results on the Hubble constant are wrong. However, we have shown that our estimates turn out to be in good agreement with all the previous ones obtained both with lensing based method and other distance ladders (such as Cepheid or SnIa). It is worth to note that the lensing observables are mainly determined by the mass distribution in the outer parts of the dark halos and this latter may be well described by models with a single power\\,-\\,law radial profile. Our models differ from the one predicted by numerical simulations only in the inner parts which are less important in determining the images positions and the time delays. It is thus expected that the difference does not introduce any serious systematic error. Anyway this does not mean that the results from numerical simulations are wrong since the statistic on the lensing systems considered is too low. To understand whether these models are really able to reproduce the lensing observables (number and position of images and time delay ratios) in quadruply imaged systems and whether they introduce a significative change in the estimate of $H_0$, one has to wait for more quadruple lenses with measured time delays. At the same time, it should be interesting to generalize our method in order to allow also a varying slope $\\alpha$ of the radial profile thus possibly leading also to some constraints from the observed quadruply imaged QSOs. Finally we note that the number of observed quadruple systems to which our method may be applied is considerable so that we have just to wait for the measure of the time delays to finally get an accurate estimate of the Hubble constant competitive with the ones coming from local estimators." + }, + "0112/astro-ph0112096_arXiv.txt": { + "abstract": "The capture rates of stars and dark particles onto supermassive black holes depend strongly on the spatial and kinematical distribution of the stellar and dark matter at the centre of bulges and elliptical galaxies. We here explore the possibility that all ellipticals/bulges have initially isothermal cusps ($\\rho \\propto r^{-2}$). If the orbits can be adequately randomized a significant fraction of the total mass of black holes in the bulges of galaxies will be due to the capture of stars and dark matter. The dark matter fraction of the total mass captured may be as high as 20--40 percent for typical cold dark matter halos. A tight relation $M_{\\bullet} \\sim 10^8 (\\sigma_*/200 \\kms)^5$ between black hole mass and stellar velocity dispersion can arise at the high mass end ($M_{\\bullet} \\ge 10^8 \\msun$) if these giant black holes grow primarily by the capture of stars without tidal disruption. For smaller black holes a shallower $M_{\\bullet}-\\sigma_*$ relation with larger scatter is expected. Efficient randomization of the orbits can be due to remnant accretion discs or the dense central regions of infalling satellites which can avoid tidal disruption and sink to the sphere of influence by dynamical friction. The presence of an isothermal cusp and the reduction of the relaxation time scale at the sphere of influence enhance the estimated tidal disruption rate of stars to $\\sim 10^{-4}$ -- $10^{-2}\\yr^{-1}$ per galaxy. Disruption flares in bright galaxies may thus be as frequent as a few percent of the supernovae rate at moderate redshifts when the galaxies still had an isothermal cusp. The efficient replenishment of the loss cone also explains why the supermassive binary black holes expected in hierarchically merging galaxies do generally coalesce as suggested by the observed relation between black hole mass and the inferred mass of stars ejected from an isothermal cusp. ", + "introduction": "The confidence in the measurement of the masses of supermassive black holes in nearby galaxies has significantly increased in the last couple of years (see Merritt \\& Ferrarese 2001 for a recent review). This is mainly due to the newly established correlation between black hole mass and velocity dispersion of the bulge of the host galaxy (Gebhardt et al. 2000, Ferrarese \\& Merritt 2000). Most if not all galactic bulges appear to contain a black hole with mass $M_{\\bullet} \\propto \\sigma_*^{4-5}$. For the published samples of reliable black hole masses this correlation appears much tighter than that between bulge mass and bulge luminosity (Kormendy \\& Richstone 1995, Magorrian et al.~1998). A number of suggestions have been made which can explain the slope of the correlation (Silk \\& Rees 1998; Haehnelt, Natarajan \\& Rees 1998; Kauffmann \\& Haehnelt 2000; Haehnelt\\& Kauffmann 2000a; Ostriker 2000; Burkert \\& Silk 2001; Adams, Graf \\& Richstone 2001) but little has been offered to explain its apparent tightness. The physical processes invoked to regulate the black hole mass depend generally on the conditions close to the black hole at radii much smaller than those at which the velocity dispersion of stars is measured. Moreover, the scatter in correlations of observed properties of galaxies like the Tully-Fisher and Faber-Jackson relation is much larger. Haehnelt\\& Kauffmann (2000b) demonstrated that their model is consistent with the observed scatter. However, the small scatter does not occur naturally in such a model where galaxies build up by hierarchical merging. Should the tightness of the correlation between black hole mass and bulge velocity dispersion stand the test of time a physical mechanism which links the black hole mass and the velocity dispersion of the stars in the bulge more directly may be required. Capture of stars by the black hole for instance is a process that depends straightforwardly on the stellar velocity dispersion (Rees 1988). The main problem is that the orbits of stars with sufficiently low angular momentum are generally assumed to be rapidly depleted, inhibiting efficient growth by accretion of stars (Sigurdsson \\& Rees 1997, Magorrian \\& Tremaine 1999, Syer \\& Ulmer 1999). The density profiles of bright ellipticals exhibit pronounced breaks within which the density profile becomes significantly shallower than isothermal at radii of a few hundred parsecs (Gebhardt 1996). Very little feeding of the black hole can come from stars at large radii, where the loss cone becomes prohibitively small. These shallower cores may however have formed very recently due to the ejection of stars by the supermassive binary black holes (Milosavljevi\\'c \\& Merritt 2001, Ravindranath, Ho \\& Filipenko 2002) expected in hierarchical merging galaxies (Kauffmann \\& Haehnelt 2000). Strong supports for this idea come from the observed correlation between the mass of the black hole and the mass inferred to be ejected if the galaxy started out with a cusp with a density profile close to isothermal. This makes it likely that bulges form with a stellar density distribution which is close to isothermal all the way down to the sphere of influence of the supermassive black hole. The number of stars on low-angular momentum orbits and thus the rate of capture of stars by the supermassive at the centre of nearby galaxies must have been much larger prior to the destruction of the isothermal cusp. Here we explore this idea in more detail. ", + "conclusions": "The capture of stars and dark matter particles from orbits with sufficiently low angular momentum to pass the event horizon contributes significantly to the total mass of black holes in the bulges of galaxies if all bulges initially had an isothermal cusp and if no depletion of these orbits occurs. The dark matter fraction of the total mass captured is 20--40 percent for typical CDM-like halos. A tight relation between black hole mass and stellar velocity dispersion of the form $M\\propto \\sigma_*^{5}$, very similar to the observed relation, arises if the black holes in the bulges of galaxies gain most of their mass by this mechanism. The relation is then expected to flatten at black holes masses smaller than $10^{8} \\msun$ where capture and tidal disruption of stars outside the horizon becomes important. Accretion of gas during the active QSO phase with its peak at redshift $z\\sim 2-3$ already accounts for a significant fraction ($\\ge 30 \\%$) of the total mass density in black holes. The capture of stars may thus tighten and refine the exact slope of a pre-existing rough correlation between black hole mass and bulge properties. It will thus be interesting to see if the tightness of the correlation persists when sample sizes get larger. In our model bright ellipticals initially had isothermal cusps and stars are predominantly captured from the sphere of influence. Relaxation of energy momentum is not required to fill the loss cone in these isothermal cusps. The randomization of stellar orbits is enhanced over the classical two-body relaxation because of processes such as the presence of long-lived remnant accretion discs left behind from phases of rapid accretion, the infall of smaller galaxies predicted by hierarchical models of galaxy formation and perhaps also the sinking of massive black holes left over from a pre-galactic episode of star formation. The presence of an isothermal cusp and the repopulation of low-angular momentum orbits will considerably increase the disruption rate of stars by supermassive black holes in the bulges of bright galaxies at moderate redshifts. Searches for high-reshift supernovae may thus detect disruption flares as frequently as a few percent of the supernovae rate. The presence of an isothermal cusp and efficient loss cone refilling will also accelerate the hardening of supermassive binary black holes so that they will generally reach the separation where they can spiral together within a Hubble time due to the emission of gravitational waves. The latter strongly supports the idea that the cores of elliptical galaxies have formed very recently due to ejection of stars by supermassive binary black holes." + }, + "0112/astro-ph0112119_arXiv.txt": { + "abstract": "The distribution of differential time delays $\\Delta t$ between images produced by strong gravitational lensing contains information on the mass distributions in the lensing objects as well as on cosmological parameters such as $H_0$. We derive an explicit expression for the conditional probability distribution function of time delays $P(\\Delta t\\,|\\,\\theta)$, given an image separation between multiple images $\\theta$, and related statistics. We consider lensing halos described by the singular isothermal sphere (SIS) approximation and by its generalization as proposed by Navarro, Frenk, \\& White (NFW) which has a density profile $\\rho \\propto r^{-\\alpha}$ in the innermost region. The time delay distribution is very sensitive to these profiles; steeper inner slopes tend to produce larger time delays. For example, if $H_0=70\\,{\\rm km\\,s^{-1}Mpc^{-1}}$, a $\\Lambda$-dominated cosmology and a source redshift $z_{\\rm S}=1.27$ are assumed, lenses with $\\theta=5^{''}$ produce a time delay of $\\Delta t[{\\rm yr}]=1.5^{+1.7}_{-0.9}$, $0.39^{+0.37}_{-0.22}$, $0.15^{+0.11}_{-0.09}$, and $0.071^{+0.054}_{-0.038}$ (50\\% confidence interval), for SIS, generalized NFW with $\\alpha=1.5$, $\\alpha=1.0$, and $\\alpha=0.5$, respectively. At a fixed image separation, the time delay is determined by the difference in the lensing potential between the position of the two images, which typically occur at different impact parameters. Although the values of $\\Delta t$ are proportional to the inverse of $H_0$, $P(\\Delta t\\,|\\,\\theta)$ is rather insensitive to all other cosmological model parameters, source redshifts, magnification biases and so on. A knowledge of $P(\\Delta t\\,|\\,\\theta)$ will also be useful in designing the observing program of future large scale synoptic variability surveys and for evaluating possible selection biases operating against large splitting lens systems. ", + "introduction": "The cold dark matter (CDM) scenario predicts relatively cuspy dark matter halos. On the basis of systematic N-body simulations, \\citet[hereafter NFW]{navarro96,navarro97} found that the density profile obeys the ``universal'' form $\\rho(r)\\propto r^{-1}(r+r_{\\rm s})^{-2}$ irrespective of the underlying cosmological parameters, the shape of the primordial fluctuation spectrum and the formation histories. Recent high resolution simulations suggest even steeper cusps $\\rho\\propto r^{-1.5}$ in the innermost region \\citep{moore99,fukushige01a,fukushige01b}, and a weak dependence of the inner slope on the halo mass is also reported \\citep{jing00a}. The statistics of strong gravitational lensing have been used to probe the density profiles of dark halos, e.g., multiple QSO images \\citep*{fox01,keeton01,wyithe01,li01,takahashi01} and the long thin arcs \\citep*{williams99,meneghetti01,molikawa01,oguri01}. These theoretical studies concluded that gravitational lensing rates are extremely sensitive to the inner slope of dark halos. The lensing statistics of small separations, however, will also be affected by gas cooling and clumpiness in the host halo \\citep{keeton98,porciani00,kochanek01,li01}. Thus multiple QSO images with intermediate or large separations, $\\theta\\gtrsim 5^{''}$, are more relevant in constraining the density profile of ``pure\" dark halos. At present, several lensing surveys at large separations have been carried out; e.g., the Jodrell-Bank VLA Astrometric Survey and the Cosmic Lens All Sky Survey \\citep[JVAS/CLASS; e.g.,][]{browne01} and Arcminute Radio Cluster-lens Search \\citep*[ARCS; e.g.,][]{phillips01}. The JVAS/CLASS sample comprises 10,499 radio sources. An explicit search for lenses has detected no lenses with image separations $6^{''}<\\theta<15^{''}$ \\citep{phillips00}, while 18 gravitational lenses with $0.3^{''}<\\theta<3^{''}$ were found \\citep{helbig00}. ARCS also produced a null result for lensing events with $15^{''}<\\theta<60^{''}$ from 1,023 extended radio sources. The lack of large separation images is only marginally consistent if the usual NFW profile and $\\Lambda$-dominated CDM model are assumed \\citep{li01,keeton01}, but it also may be ascribed in part to an effect of the longer time delays between more widely separated images. Any intrinsic variability of the QSO will result in images which less resemble each other for longer delays \\citep{phillips01}. So far, time delays between two images have been used primarily to estimate the Hubble constant $H_0$ using a detailed lens model of each system \\citep[e.g.,][]{grogin96,barkana99}. But the importance of mass distribution in estimating Hubble constant has been recognized \\citep{impey98,keeton00,witt00} and this led to an attempt to constrain the galaxy mass profile from time delays using Monte Carlo simulations \\citep{rusin00}. Instead, in this paper, we consider the statistics of the time delay effect analytically. We derive an expression for the cumulative joint probability, i.e., the probability that the time delay is larger than $\\Delta t$ and the image separation is $\\theta$, $P(>\\!\\Delta t, \\theta)$, and also various related statistics. They allow us to estimate the range of probable time delays for a given image separation and the extent to which the intrinsic time variability of quasars affects observed strong gravitational lensing rates. Our most important result is that the distribution of delays is quite sensitive to the density profiles of the lensing objects. For example, an NFW density profile predicts median time delays a factor of three or so smaller than the density profile proposed by \\citet{moore99} and \\citet{fukushige01a,fukushige01b} for the same $H_0$ value. While the lenses for which time delays are currently observed are dominated by barionic component, we can constrain the density profile of dark halos if a sample of time delays with large separations becomes available. It turns out that the time delay statistics are fairly insensitive to other uncertainties, such as the magnification bias and various cosmological parameters and therefore become a relatively reliable estimator for the density profile of dark halos. Of course, all delay values are linearly proportional to the inverse Hubble constant, and we here assume its value to be $H_0=70\\,{\\rm km\\,s^{-1}Mpc^{-1}}$, which is consistent with the final result of Hubble Space Telescope Key Project \\citep{freedman01}, throughout the remainder of our discussion. The outline of this paper is as follows. In \\S 2, we briefly describe the usual formulation of gravitational lensing statistics. Section 3 presents the analytic formulation of time delay statistics. Our main results are shown in \\S 4. Finally we summarize the main results and discuss their application in \\S 5. ", + "conclusions": "In this paper, we formulated the differential time delay probability distribution and examined its dependence on lens density profiles as well as on magnification bias, cosmological model, and source redshift. We found that the probability distribution of time delays depends most sensitively on the inner slope of the lens density profiles. The difference between the various density profiles we examined is very large, more than one order. For example, if $H_0=70\\,{\\rm km/s/Mpc}$, $\\Lambda$-dominated cosmology and a source redshift $z_{\\rm S}=1.27$ are assumed, lenses with $\\theta=5^{''}$ induce the time delay; $\\Delta t[{\\rm yr}]=1.5^{+1.7}_{-0.9}$, $0.39^{+0.37}_{-0.22}$, $0.15^{+0.11}_{-0.09}$, and $0.071^{+0.054}_{-0.038}$ (50\\% level), for SIS, generalized NFW with $\\alpha=1.5$, $\\alpha=1.0$, and $\\alpha=0.5$, respectively. On the other hand, the effects of cosmological models, source redshifts or magnification bias are rather small, aside from the well known and important linear dependence on the inverse Hubble constant. Moreover, the delay distributions are quite insensitive to the normalization of the overall lensing rate, because the conditional probability is defined by the joint probability divided by the usual lensing probability and the differing of normalizations are almost canceled out. Therefore, one could strongly constrain the core structure of dark halos if a large sample of time delays becomes available in future systematic survey. Although existing lens systems are individually modeled in detail including the central density profiles, such careful treatment of each lens system in upcoming very large samples may not always be practical. Thus a {\\it statistical} treatment of time delays is informative. For example, we can predict the range of probable time delays of large separation lenses from this formulation if the density profile of dark halos is fully settled. This would allow estimation of a plausible range of time delays for {\\it future} lens systems even when the lensing object has not been identified. Comparison with the meager existing observational data suggests that the density profile has a rather steep cusp, $\\alpha\\sim2$, although it is already known that the current observed systems are dominated by baryonic component. Moreover, it seems that the smaller separation lenses prefer steeper inner profiles. This result may support a two-population model, i.e., galactic mass halos with a steep inner slope ($\\alpha\\sim2$) and cluster mass halos with a shallower inner slope ($\\alpha\\lesssim1.5$), as proposed by \\citet{keeton98} (see also \\citealp{porciani00,li01}). To see whether there are two (or more) populations of halos, time delay data for various separations, especially separations with $5^{''}\\lesssim\\theta\\lesssim10^{''}$, will be essential. Another application of time delay statistics is that future large samples of lenses will probably be systematically monitored by one or more of the ambitious synoptic surveys now being planned. Designing an efficient sampling rate and observing strategy requires an idea of the range of time delays that might reasonably be expected." + }, + "0112/astro-ph0112433_arXiv.txt": { + "abstract": "The Two Micron All Sky Survey (2MASS) is finding previously unidentified, luminous red active galactic nuclei (AGN). This new sample has a space density similar to, or greater than, previously known AGN, suggesting that a large fraction of the overall population has been missed. {\\it Chandra\\/} observations of a well-defined subset of these objects reveal that all are X-ray faint, with the reddest sources being the faintest in X-rays. The X-ray hardness ratios cover a wide range, generally indicating \\nh $\\sim 10^{21-23}$ cm$^{-2}$, but the softest sources show no spectral evidence for intrinsic absorption. These characteristics suggest that a mix of absorbed, direct emission and unabsorbed, scattered and/or extended emission contributes to the X-ray flux, although we cannot rule out the possibility that they are intrinsically X-ray weak. % This population of X-ray faint, predominantly broad-line objects could provide the missing population of X-ray absorbed AGN required by current models of the cosmic X-ray background. The existence of AGN which display both broad emission lines and absorbed X-rays has important implications for unification schemes and emphasizes the need for care in assigning classifications to individual AGN. ", + "introduction": "The realization that obscuration plays a critical role in the classification of AGN inspired a fundamental change in our understanding of the phenomenon. Not only does the ``Unified Scheme'', in which narrow emission line (type 2) AGN are interpreted as edge-on broad emission line (type 1) AGN, provide a basis for new observations and theoretical models, we also realize that many AGN may be nearly invisible in UV-excess surveys ({\\it e.g.,} Webster \\etal\\ 1995; Masci \\etal\\ 1999). This expanded idea of what comprises an AGN means that their current number density may be significantly underestimated ({\\it e.g.,} Sanders and Mirabel 1996). Ramifications include revisions of the fraction and types of galaxies that harbor an active nucleus, the energy density of ionizing flux in the young universe, and the nature of the X-ray and far-IR backgrounds. Although {\\it IRAS\\/} provided the first significant sample of extragalactic objects in which the bulk of the luminosity emerges as reprocessed radiation in the IR (Soifer \\etal\\ 1984), its sensitivity was sufficient to catalog only the most nearby and/or luminous AGN. The Two Micron All Sky Survey (2MASS) is yielding a much deeper catalog of near IR-selected AGN (Cutri \\etal\\ 2001) by selecting sources with $J-K_s > 2$ from the high galactic latitude 2MASS Point Source Catalog. Spectroscopic follow-up of red candidates reveals $\\sim$75\\% are {\\it previously-unidentified} emission-line AGN, with $\\sim$80\\% of these showing broad optical emission lines (Type 1: Seyfert 1 and QSO), and the remainder being narrow-line objects (Type 2: Seyfert 2, QSO 2, and LINER; Cutri \\etal\\ 2001). They span a redshift range $0.1$25,000 such objects over the sky. The objects have unusually high optical polarization levels, with $\\sim$10\\% showing $P>3$\\% indicating a significant contribution from scattered light (Smith \\etal\\ 2001). {\\it ROSAT\\/} found that, while known AGN dominate the soft ($0.1-2.0$~keV) cosmic X-ray background (CXRB; Lehmann \\etal\\ 2000), an additional population of heavily absorbed AGN would be required to account for the harder high-energy spectrum (Comastri \\etal\\ 1995). To match both the CXRB spectrum and the observed hard X-ray number counts of pre-{\\it Chandra\\/} surveys (Fiore \\etal\\ 2001), the X-ray absorbed AGN population is estimated to outnumber unabsorbed AGN by $\\sim$4:1 and perhaps to increase with $z$ (Gilli \\etal\\ 2001; Comastri \\etal\\ 2001). Although the ratio of Type 2 to Type 1 AGN in the local universe is consistent with this: $\\sim$$2-4$ (Maiolino \\& Rieke 1995; Huchra \\& Burg 1992), a dominant population of X-ray absorbed AGN at z\\gax 0.1 has yet to be found. Possible identifications include ADAF (advection dominated accretion flow) galaxies (Di Matteo \\etal\\ 1999) and narrow emission-line X-ray galaxies with flat/absorbed X-ray spectra, luminous IR galaxies (Risaliti \\etal\\ 2000), and a subset of broad-line AGN that are X-ray absorbed. The {\\it Chandra\\/} X-ray Observatory (Weisskopf \\etal\\ 2000) and {\\it XMM-Newton\\/} (Jansen \\etal\\ 2001), with their faint flux limits and broad energy sensitivity ($\\sim$$0.5-10$ keV), are finding objects in sufficient numbers to explain $60-80$\\% of the CXRB, including a significant number of hard spectrum sources. These correspond to both optically faint objects and bright, nearby, but otherwise normal elliptical galaxies (Hornschemeir \\etal\\ 2000; Barger \\etal\\ 2001; Giacconi \\etal\\ 2001), as well as more traditional, broad-line AGN. Near-IR observations of the latter reveal featureless, red continua ($1.5 < J-K_S < 2.5$) which, combined with their optical colors, are consistent with the moderate amounts of absorption required to match the CXRB (Compton-thin; equivalent neutral hydrogen column density $N_{\\rm H}~<~10^{24}$~cm$^{-2}$; Crawford \\etal\\ 2001; Barger \\etal\\ 2001). Thus, evidence is mounting that absorbed AGN are indeed important contributors to the CXRB. The area covered by the deepest {\\it Chandra\\/} surveys is small ($\\sim$0.16 sq. deg.), resulting in large statistical uncertainties, particularly at low $z$. Given the high surface density and similarities of the new population of 2MASS AGN to the smaller {\\it Chandra\\/} sample, a census of the X-ray properties of the 2MASS AGN will likely be an essential ingredient in their understanding and will yield an estimate of whether this previously missed population is sufficient, alone, to explain the shortfall in the CXRB. ", + "conclusions": "" + }, + "0112/astro-ph0112009_arXiv.txt": { + "abstract": "Archeops is a balloon-borne experiment designed to measure the temperature fluctuations of the CMB on a large region of the sky ($\\simeq 30\\%$) with a high angular resolution (10 arcminutes) and a high sensitivity ($60\\mu\\mathrm{K}$ per pixel). Archeops will perform a measurement of the CMB anisotropies power spectrum from large angular scales ($\\ell\\simeq 30$) to small angular scales ($\\ell \\simeq 800$). Archeops flew for the first time for a test flight in July 1999 from Sicily to Spain and the first scientific flight took place from Sweden to Russia in January 2001. The data analysis is on its way and I present here preliminary results, realistic simulations showing the expected accuracy on the measurement of the power spectrum and perspectives for the incoming flights (Winter 2001/2003). \\vspace{1pc} ", + "introduction": "\\subsection{CMB physics} In the framework of Big-Bang theory, the Universe started with a hot and dense phase about 15 billion years ago and cooled down while expanding. The first neutral atoms formed when the temperature was about 13.6 eV (160000 K), but due to the large number of photons compared to baryons (ratio $\\simeq 10^9$), the Universe remained ionized until the temperature dropped below 0.3 eV (3000 K). At this moment, the mean free path of the photons increased drastically so that the photons that scattered at this time have not interacted with matter since then. This moment is known as {\\em matter-radiation decoupling} or {\\em recombination}. Those photons cooled down with the expansion of the Universe and are know observed at a temperature of 2.7 K. As the matter and radiation were at thermal equilibrium before decoupling, these photons have a pure blackbody spectrum and are homogeneously distributed on the celestial sphere. This radiation is known as the {\\em Cosmic Microwave Background} (hereafter CMB). The discovery of the CMB by Penzias and Wilson~\\cite{penzias_wilson} and its interpretation in terms of a Big-Bang relic by Dicke and collaborators~\\cite{dicke} was a major argument for the Big-Bang theory~\\cite{gamow,alpher_herman}. The CMB temperature was measured to be highly isotropic but tiny anisotropies were expected. These temperature fluctuations reflect the density fluctuations on the last scattering surface. These are necessary to explain the presence of structures in the Universe such as galaxies and clusters. The CMB anisotropies were discovered by the COBE satellite with a {\\em rms} amplitude of about 30 $\\mu\\mathrm{K}$~\\cite{smoot} at scales larger than 7 degrees. COBE also measured its spectrum with high precision~\\cite{mather,fixsen} proving its pure blackbody nature. The CMB anisotropy typical physical size in the last scattering surface can be theoretically predicted while its angular size as seen from here and now depends on the geometry of the Universe along the path of the photons. Hence, mapping the CMB anisotropies is a powerful cosmological test. The two competing paradigms for the origin of structures in the Universe, namely inflation and topological defects, predict significantly different distributions for the former density fluctuations. These distributions propagate to us in a cosmological parameters dependent way to describe the temperature anisotropies that we expect on the sky\\footnote{Freely available numerical codes, such as CMBFast~\\cite{cmbfast}, have been developed for this purpose.}. It is therefore of deep interest to investigate their angular distribution and compare the measurements to cosmological models. The temperature anisotropies on the sky are commonly described via their spherical harmonics expansion, \\begin{equation} \\frac{\\delta T}{T}\\left(\\theta,\\phi\\right)=\\sum_{\\ell=0}^\\infty\\sum_{m=-\\ell}^\\ell a_{\\ell m}Y_{\\ell m}\\left(\\theta,\\phi\\right), \\end{equation} where $\\ell$ is the multipole index, inversely proportional to the angular scale (1 degree roughly corresponds to $\\ell=200$). The angular power spectrum of the temperature fluctuations of the CMB is defined as: \\begin{equation} C_\\ell=\\frac{1}{2\\ell+1}\\sum_{m=-\\ell}^\\ell\\left|a_{\\ell m}\\right|^2. \\end{equation} The evolution of the angular power spectrum of the CMB as a function of $\\ell$ can be splitted into three major regions (see Figure~\\ref{speccmb}): \\begin{itemize} \\item On the low-$\\ell$ part (large angular scales) no particular structure is expected as we are considering physical sizes on the last scattering surface larger than the horizon at the epoch of decoupling. No physical process is expected to have modified those fluctuations since the early Universe. \\item Between $\\ell\\simeq 30$ and $\\ell\\simeq 1000$ (degree and sub-degree scales) we a re considering structures that had time to collapse and experience acoustic oscillations between the matter-radiation equality and the matter radiation decoupling. We therefore expect a series of acoustic peaks (the first one being located around $\\ell=200$, corresponding to the size of the horizon at the epoch of decoupling) in the case of inflationary-like early Universe models where the oscillations are in phase. In the case of isocurvature fluctuations (such as topological defects), the oscillations are not in phase and a large bump is expected, but no multiple peaks. \\item In the large $\\ell$ part (arcminute scales and below), the power is expected to drop drastically due to the finite thickness of the last scattering surface and to the finite value of the mean free path of the photons before decoupling. \\end{itemize} \\begin{figure}[htb] \\resizebox{\\hsize}{!}{\\includegraphics{cl_2001.ps}} \\caption{Expected CMB power spectrum $\\left(\\sqrt{\\frac{\\ell\\left(\\ell+1\\right)C_\\ell}{2\\pi}}\\right)$ for inflationary-like primordial density fluctuations (black curves) for three different cosmological models along with the latest measurements from BOOMERanG, MAXIMA and DASI and the earlier measurements from COBE.} \\label{speccmb} \\end{figure} \\subsection{Recent results} Our knowledge of the CMB power spectrum has been significantly improved since COBE measurements~\\cite{smoot} by two balloon-borne experiments, BOOMERanG~\\cite{netterfield} and MAXIMA~\\cite{hanany}, and a ground based interferometric experiment, DASI~\\cite{halverson}. These results are shown in Figure~\\ref{speccmb}. We observe that the very low part of the power spectrum is highly constrained by the COBE points while the high-$\\ell$ acoustic region is constrained by the three recent experiments showing undoubtly the multiple peak feature that is expected from inflation. This set of data therefore strongly disfavors topological defects as seeds for the structure formation in the Universe. Comparing this set of data with theoretical power spectra have lead to the estimation of the cosmological parameters~\\cite{netterfield,balbi,bond}. The favored model is dominated by dark energy (around 70\\%) such as cosmological constant or quint essence (in agreement with high redshift type Ia supernovae measurements~\\cite{perlmutter,schmidt}). The 30\\% of matter consists mainly of dark matter and the amount of baryons is in agreement with Big-Bang nucleosynthesis predictions and light elements abundances measurements. One of the great news coming from these results is that they agree very well with other measurements of the cosmological parameters obtained with completely different methods: large scale structure observation, lensing, type Ia supernovae and light elements abundances. We are now heading towards a concordance model. \\subsection{Motivations for an intermediate scale experiment} There is a part of the CMB power spectrum that lacks measurements: between the large angular scales measurements from COBE (around $\\ell=20$) and the largest angular scales from BOOMERanG, MAXIMA and DASI (above $\\ell\\simeq 75$). This can be easily understood as when you are trying to measure large $\\ell$, you concentrate on a small patch of the sky in order to reach high signal to noise ratio. This implies loosing all information on large scales. On the other hand, COBE covered the entire celestial sphere but with a poor 7 degrees resolution that limited the measurements to the low $\\ell$. The intermediate part of the power spectrum, despite being difficult to measure, is of high interest.: first, this is the part of the power spectrum that is the most sensitive to the early Universe physics. Second, linking COBE measurements to high-$\\ell$ measurements with one single experiment would ensure that no calibration problems are affecting the data. Archeops is a balloon-borne experiment designed to measure the CMB temperature angular power spectrum from the large multi-degree scales ($\\ell\\simeq 30$) to the small sub-degree scales ($\\ell\\simeq 800$). This is achieved via a large sky coverage (around 35\\%) along with a high angular resolution (10 arcminutes). Archeops is also very similar to Planck High Frequency Instrument and can therefore be considered as a real size testbed for Planck-HFI. ", + "conclusions": "" + }, + "0112/astro-ph0112523_arXiv.txt": { + "abstract": "\\input abstract.tex ", + "introduction": " ", + "conclusions": "" + }, + "0112/astro-ph0112379_arXiv.txt": { + "abstract": "{\\normalsize We show that supernova neutrinos can be studied by observing their charged-current interactions with $^{100}$Mo, which has strong spin-isospin giant resonances. Information about both the effective temperature of the electron-neutrino sphere and the oscillation into electron neutrinos of other flavors can be extracted from the electron (inverse $\\beta$) spectrum. We use measured hadronic charge-exchange spectra and the Quasiparticle Random Phase Approximation to calculate the charged-current response of $^{100}$Mo to electron neutrinos from supernovae, with and without the assumption of oscillations. A scaled up version of the MOON detector for $\\beta \\beta$ and solar-neutrino studies could potentially be useful for spectroscopic studies of supernova neutrinos as well.}\\\\ \\vspace{.1cm} \\noindent PACS : 23.40-s,14.60.Pq, 26.65.+t, 95.55.Vj ", + "introduction": " ", + "conclusions": "" + }, + "0112/astro-ph0112465_arXiv.txt": { + "abstract": "We report long-term simultaneous optical and (RXTE) X-ray observations of the soft X-ray transient and low mass X-ray binary X1608$-$52 spanning from 1999 to 2001. In addition to the usual X-ray outburst and quiescent states, X1608$-$52 also exhibits an extended low intensity state during which the optical counterpart, QX~Nor, is found to be about two magnitudes brighter than during quiescence. We detect optical photometric variability on a possible period of 0.5370~days with a semi-amplitude of $\\sim 0.27$ mag in the $I$ band. The modulation could be orbital but is also consistent with a scenario invoking a superhump with decreasing period. We discuss the possibilities to distinguish between the orbital and superhump period cases. Observations of QX~Nor during quiescence indicate an F to G type main sequence secondary while theoretical considerations argue for an evolved mass donor. Only an evolved mass donor would satisfy the condition for the occurrence of superhumps. ", + "introduction": "Soft X-ray transients (SXTs) are a subclass of the low mass X-ray binaries (LMXBs), mass-transfer systems where matter donated by a late-type star is accreted onto either a neutron star (NS) or a black hole. SXTs are characterized by prominent outbursts separated by long periods of quiescence. The physical mechanism underlying the outbursts remains under debate, although the favored scenario at present is some variation of the disk instability model (see the recent review by Lasota 2001). During an outburst, the luminosity of SXTs increases by several orders of magnitude at all wavelengths. While the optical light is dominated by the accretion disk and the heated face of the companion during the outburst as in the persistent LMXBs, the quiescent state of SXTs offers the rare opportunity to study the mass donor ``uncontaminated'' by other sources of emission. Observations of the mass donor in quiescence can provide the orbital period of the system through ellipsoidal variations, the spectral type and radial velocity curve of the donor star, and also constrain the mass of the compact object. The majority of SXTs appear to contain black hole primaries and this group provides some of the best black hole candidates with the most stringent mass constraints. X1608$-$52 is a member of the rarer group of SXTs with neutron star primaries (based on the detection of type I X-ray bursts). X-ray outbursts from this system have been recorded since the early 1970's. While some of the data suggested a recurrence pattern of $\\sim 600$ days (Murakami et al.\\ 1980), Lochner \\& Roussel-Dupre (1994) reported a more complex outburst history. Hasinger \\& van der Klis (1989) classified X1608$-$52 as an Atoll source. Kilohertz quasi-periodic oscillations (QPOs) were discovered during the decay of the 1996 outburst (van Paradijs et al.\\ 1996, Berger et al.\\ 1996), twin QPO peaks with varying separation during the 1998 outburst (Mendez et al.\\ 1998), and finally possible sideband oscillations (Jonker, Mendez \\& van der Klis 2000). The optical counterpart of X1608$-$52, QX~Normae, was discovered by Grindlay \\& Liller (1978) after an outburst in 1977. Wachter (1997) detected QX~Nor in the IR and optical a few months after its outburst in 1996. If X1608$-$52 was to adhere to the recurrence pattern of its last two outbursts (February 1996 and 1998), the next outburst was expected to occur early in 2000. We report here on long-term optical and X-ray observations spanning from 1999 to 2001. ", + "conclusions": "We have identified three distinct states in the X-ray and optical light curves of X1608$-$52. In addition to the previously recognized outburst and (true) quiescence states, X1608$-$52 also exhibits an extended low intensity state during which the optical counterpart, QX~Nor, is found to be about two magnitudes brighter than during true quiescence. The LIS is characterized by the presence of a luminous accretion disk in the system. The optical counterpart displays pronounced variability that can be attributed to either the varying aspect of the X-ray irradiated face of the mass donor with an orbital period of 0.5370~days, or a time variable superhump period with $P_0=0.554$~days and $\\dot{P}=-0.00024$. This ambiguity could be resolved with photometric monitoring during a future bright state of the system. If the mass donor is assumed to be a main sequence star, an F to G type star is consistent with the current observations. However, theoretical considerations argue for an evolved mass donor with a mass of $\\sim 0.32 M_\\odot$ (for a $1.4 M_\\odot$ NS primary). Only an evolved mass donor would satisfy the condition for the occurrence of superhumps." + }, + "0112/astro-ph0112186_arXiv.txt": { + "abstract": "{We study photometric variability among the optical counterparts of X-ray sources in the old open cluster M\\,67. The two puzzling binaries below the giant branch are both variables: for S\\,1113 the photometric period is compatible with the orbital period, S\\,1063 either varies on a period longer than the orbital period, or does not vary periodically. For the spectroscopic binaries S\\,999, S\\,1070 and S\\,1077 the photometric and orbital periods are similar. Another new periodic variable is the main-sequence star S\\,1112, not known to be a binary. An increase of the photometric period in the W\\,UMa system S\\,1282 (AH\\,Cnc) is in agreement with a previously reported trend. Six of the eight variables we detected are binaries with orbital periods of 10 days or less and equal photometric and orbital periods. This confirms the interpretation that their X-ray emission arises in the coronae of tidally locked magnetically active stars. No variability was found for the binaries with orbital periods longer than 40 days; their X-ray emission remains to be explained. ", + "introduction": "Twenty five members of the old open cluster M\\,67 have been detected in X-rays (Belloni et al. 1998). At the age of M\\,67 (4 Gyr, Pols et al. 1998) the rotation of single stars is too slow to generate detectable X-rays. Therefore, the X-ray emission of many M\\,67 sources probably arises in interacting binaries. Indeed, one source is known to be a cataclysmic variable. Nine sources are binaries with orbital periods of 10 days or less, presumably RS\\,CVn type systems, whose X-rays are due to the coronae of magnetically active stars forced to corotate by tidal interaction (see Table~\\ref{Xlist}). However, not all X-ray sources are binaries, e.g. one source is a hot white dwarf, and some others are stars which do not show signs of binarity. There are five peculiar binaries (S\\,1040, S\\,1063, S\\,1072, S\\,1082 and S\\,1113) whose evolutionary statuses we currently do not understand. They are found in the colour-magnitude diagram in locations which cannot be reproduced by combining the light from any two members on the main sequence, subgiant and/or giant branches. Two of them have long orbital periods which exclude strong tidal interaction. A sixth binary (S\\,1237) with a long orbital period lies to the blue of the giant branch which can be explained by the superposition of the light of a giant and a turnoff star. A spectroscopic study of these peculiar systems was presented in van den Berg et al. (1999); see their Fig.~1 for the position of these systems in the colour-magnitude diagram of M\\,67. Here we report our photometric study of these sources and of other X-ray sources that happen to be in the same fields of view. The photometry of S\\,1082 is published separately (van den Berg et al. 2001). The observations and analyses are described in Sect.~\\ref{obs}. Results are presented in Sect.~\\ref{results}, followed by the interpretation and discussion -- including comparison with earlier work -- in Sect.~\\ref{disc}. Sect.~\\ref{concl} summarizes our conclusions. The variability of stars not detected in X-rays but included in our observations will be the subject of Paper II (Stassun et al. 2001). \\nocite{bellea} \\nocite{vdbergea} \\nocite{stasead} \\nocite{polsea} \\nocite{san} \\nocite{montea} \\nocite{mathlathea} ", + "conclusions": "\\label{concl} Of the twenty two X-ray sources in M\\,67 that we discuss, sixteen are spectroscopic binaries with known orbital periods. Our survey for optical photometric variables among these X-ray sources has established eight variables. Seven of these are among the sixteen binaries, the binary status of the eighth, S\\,1112, is not yet known. In addition, Gilliland et al. (1991) observed periodic optical variation in three more of the X-ray binaries with amplitudes too low to be detected by us: S\\,1019 (semi-amplitude 0.015 mag), S\\,1242 (semi-amplitude 0.0025 mag) and S\\,1040 (semi-amplitude 0.012) mag. Thus ten of the sixteen X-ray binaries in M\\,67 are optical variables at the $\\gtrsim$ 0.01 mag level. Belloni et al. (1998) have suggested that rapid stellar rotation resulting from tidal locking results in enhanced magnetic activity and X-ray emission. Fig.~\\ref{sb} shows the visual magnitude versus orbital period of the spectroscopic binaries. With the exception of S\\,1040 and S\\,1112, all variables have orbital periods less than 20 days and $V>$\\,15. In all cases but S\\,1019 and S\\,1063, the photometric period is equal to the orbital period or, in the case of S\\,1242, the orbital period near periastron. Evidently tidal locking has been established, leading to rotation of at least the primary star that is more rapid than typical for solar-mass stars at 4 Gyr. Thus our results establish a key premise of the Belloni et al. (1998) picture for the X-ray emission. Furthermore, if the cause of the observed optical variability is indeed spot modulation of the observed flux, then the presence of the required large spots is consistent with enhanced magnetic activity in these stars. The X-ray emission and optical variability properties of S\\,1019 and S\\,1063 require further investigation. Three binaries were not detected as variables despite their short orbital periods. S\\,972 is the faintest of the binary sample at $V$\\,=\\,15.37, and so its variability may have gone undetected. The X-ray luminosities of S\\,1045 ($P_{{\\rm orb}}$\\,=\\,7.6 days) and S\\,1234 ($P_{{\\rm orb}}$\\,=\\,4.3 days) are among the lowest of the binary X-ray sources and indicate low activity levels; this can explain the absence of optical variability due to spots. Rajamohan et al. (1998) have noted S\\,1234 as a possible optical variable (semi-amplitude $\\sim$0.16 mag) which suggests that time-variability of the spot phenomenon can also explain the absence of optical variation. The interpretation of S\\,1040 and of the remaining three X-ray binaries S\\,760, S\\,1072 and S\\,1237 may be the most challenging. All have long orbital periods. Given their wider separations tidal locking is not expected and so the consequent stellar rotations may be characteristic of single stars. As such, their lack of large spots and consequent photometric variability is not a surprise. Nonetheless, these binaries are X-ray sources. Their X-ray emission remains to be explained. \\begin{figure} \\centerline{\\psfig{figure=sb3a.ps,angle=-90,width=\\columnwidth,clip=t} {\\hfill}} \\caption{Visual magnitude versus orbital period of the M\\,67 binaries detected in X-rays. The size of the symbol is a measure for the logarithm of the X-ray luminosity (0.1-2.4 keV, in erg s$^{-1}$) as indicated in the figure. Eccentric binaries are indicated with trangles, binaries with eccentricities compatible with zero (within the 3$\\sigma$ error) with circles. Filled symbols are systems for which we detected photometric variability. S\\,1112 is indicated with a filled square. \\label{sb}} \\end{figure} No large radial-velocity variations were found for S\\,775 and S\\,1270 ($\\sigma$ is 0.9 km s$^{-1}$ in 12 observations spanning 5200 days and 0.7 km s$^{-1}$ in 7 observations spanning 800 days, respectively, see Mathieu et al. 1986); if these stars are binaries their periods must be relatively long. Thus their X-ray luminosities, as those of S\\,364, S\\,628 and S\\,1027 for which no radial-velocity information is available, remain unexplained. \\nocite{rajaea} \\nocite{kara} \\nocite{mathea86}" + }, + "0112/astro-ph0112135_arXiv.txt": { + "abstract": "{ We present a list of \\zae\\ ($\\Delta v < 3000$\\,\\kms) DLAs discovered during the CORALS survey for absorbers in a radio-selected QSO sample. On the assumption that \\zae\\ DLAs are neither ejecta from the QSO, nor part of the host galaxy itself, we use the number density statistics of these DLAs to investigate galaxy clustering near the QSO redshift. We find that $n(z)$ of \\zae\\ DLAs in our radio-selected QSO sample is $\\sim 4$ times larger than the number density of intervening DLAs, implying an excess of galaxies near the QSO. This result is further supported with the inclusion of the radio-loud QSOs in the FIRST survey, although the total number of DLAs is still small (4) and the result is currently only significant at approximately the $2\\sigma$ level. Since all of the \\zae\\ DLAs we identify in CORALS are found towards optically bright ($B < 20$) QSOs, there is no strong evidence (based on these limited statistics) that this population suffers from a severe dust bias. We compare our results with those from an optically-selected, radio-quiet QSO sample in order to determine whether there is evidence for an excess of galaxies around radio-loud versus radio-quiet QSOs. We find that the $n(z)$ of \\zae\\ DLAs towards radio-quiet QSOs is in agreement with the number density of intervening absorbers. This result, although currently limited by the small number statistics of our survey, supports the conclusion that radio-loud quasars are found preferentially in rich galaxy environments at high redshift. Finally, we report that one of the new \\zae\\ DLAs discovered by CORALS has some residual flux in the base of the \\lya\\ trough which may be due to \\lya\\ emission, either from star formation in the DLA galaxy or from gas photoionised by the QSO. ", + "introduction": "The formation of quasars is believed to represent an integral stage in the process of galaxy formation and evolution. The high incidence of galaxy-QSO companions and apparent mergers supports the view that quasar activity is fuelled by continued accretion of material onto the central black hole from the intercluster medium or nearby galaxies (see the review by Barnes \\& Hernquist 1992). Many of the outstanding questions in theories of the role of AGN in galaxy evolution involve the link between radio power and the properties of the QSO environment. The long-standing view is that both the morphology of the host galaxy (e.g. Hamilton, Casertano \\& Turnshek 2001), and the richness of the surrounding environment (Yee \\& Green 1987; Hall \\& Green 1998; Hall, Green \\& Cohen 1998; Hutchings et al. 1999; Teplitz, McLean \\& Malkan 1999; Cimatti et al. 2000) are different between radio-quiet QSOs (RQQs) and radio-loud QSOs (RLQs). Specifically, Yee \\& Green (1987) pointed out that the former are rarely found in clusters, while about 35\\% of RLQs show an excess of faint galaxies in their vicinity, consistent with the presence of a cluster of galaxies at the QSO redshift (see also Sanchez \\& Gonzalez-Serrano 1999). Radio galaxies, which are intimately connected with RLQs in unification models, also appear to reside in cluster environments at high redshift (Pentericci et al. 2000). However, the dichotomy between radio-loud and radio-quiet properties is now being questioned by studies which have found that almost all QSO hosts are elliptical galaxies, regardless of radio-loudness (McLure et al. 1999; Dunlop et al. 2002), and that the environments of RLQs are statistically indistinguishable from those of RQQs, at least at $z \\simlt 1$ (Hutchings, Crampton \\& Johnson 1995; McLure \\& Dunlop 2001; Wold et al. 2001; Finn, Impey \\& Hooper 2001). At higher redshifts, there is evidence that galaxy density maybe higher around RLQs, although a milder overdensity is seen towards RQQs (see review by Hutchings 2001). This ongoing discussion is made even more relevant by (a) our lack of understanding of the mechanism responsible for the radio power of RLQs and (b) the recent paradigm shift away from the long held belief that the distribution of radio luminosities is bimodal (Kellermann et al 1989; White et al. 2000; Brinkmann et al. 2000; Lacy et al. 2001). It is of primary importance to establish the extent to which radio power depends on the QSO environment, since this relationship clearly holds important clues to the physics of AGN and the role of quasars in galaxy formation. In this paper we explore the possibility of using quasar absorption lines as an unbiased probe of the high redshift QSO environment. This is a potentially powerful technique since studies of galaxy clustering around QSOs can be extended to much higher redshifts, although the obvious disadvantage is the uni-dimensional information content. Therefore, a large number of QSOs needs to be studied in order for the statistical trends of \\zae\\ absorbers to provide meaningful results. In this respect, the results presented here for the 66 QSOs of the Complete Optical and Radio Absorption Line System survey (CORALS, Ellison et al. 2002) should be regarded as preliminary, although they underline the potential of this approach, so far relatively unexploited. There is already some evidence that the posited correlation between galaxy concentration and radio-loudness may extend to QSO absorption line systems, based on an excess of associated C~IV systems (Foltz et al. 1986; Anderson et al. 1987; Foltz et al. 1988) in RLQ spectra out to 3000 \\kms\\ from the QSO redshift. However, the origin of these absorbers is unclear. The excess seems to be restricted to steep-spectrum RLQs, and may thus be related to some orientation effect (Barthel et al. 1997; Richards et al. 1999; Richards 2001). On the other hand, Richards et al. (2001) find no correlation of this excess with other orientation measures such as core-to-lobe ratio. Nevertheless, the suspicion remains that these unusually strong and highly ionised systems are probably intrinsic to the QSO and may not be providing us with specific clues about the surrounding galaxy environment (Baker et al 2001). A more fruitful line of investigation is to focus on \\zae\\ absorbers which are analogous to galaxy-scale absorbers at intervening redshifts, namely the Damped Lyman Alpha systems (DLAs). In the last 15 years, several surveys have mined the sky for DLAs, with the objective of understanding more about the high redshift galaxy population that these absorption systems are thought to represent (Wolfe et al. 1986; Lanzetta et al. 1991, 1995; Wolfe et al. 1995; Storrie-Lombardi \\& Wolfe 2000; P\\'{e}roux et al. 2001). Almost without exception, these surveys have adopted the early strategy and definitions of Wolfe et al. (1986) who included in their sample only DLAs with $N$(H~I) $\\ge 2 \\times 10^{20}$ atoms \\cm\\ that lie at absorption redshifts corresponding to $\\Delta v >$ 3000 \\kms\\ from the QSO emission redshift. The reason for imposing this velocity cut is primarily due to the unknown nature of these proximate DLAs (PDLAs), which may conceivably be caused by QSO ejecta or by absorption in the host galaxy itself. In addition, since one of the main objectives of DLA surveys has been to compile a census of H~I in the universe, proximate systems have been excluded due to the possible effect of the QSO's ionizing radiation on the absorber. For example, the intense local radiation field of the QSO is already well known to affect the distribution of the lower column density \\lya\\ forest clouds (e.g. Murdoch et al. 1986). A similar phenomenon has also been noted to affect the extended \\lya\\ haloes of low redshift galaxies (Pascarelle et al. 2001). Therefore, in the quest to gain an insight into the `normal and representative' galaxy population at high redshift, PDLAs at \\zae\\ have been almost universally discarded from the statistics of previous surveys. M\\o ller, Warren \\& Fynbo (1998), motivated by observations of \\lya\\ emission in associated DLAs, have explored the possible nature of these systems. They concluded that PDLAs are unlikely to be due to QSO ejecta, based on their lower metallicities, lack of strong high ionization lines and simpler line profiles compared to intrinsic BAL features (see Barlow et al. 1997 for a summary of intrinsic absorber signatures). Although recent models for the structure of AGN (e.g. Elvis 2000) include relatively high column density, wind-ejected clouds, the consequent UV absorbers are quite different from PDLAs both in terms of $N$(H~I) and ejection velocity (e.g. Monier et al. 2001). For example, narrow intrinsic C~IV systems, identified by means of time variability or partial coverage have relatively low H~I column densities (e.g. Petitjean, Rauch \\& Carswell 1994; Hamann, Barlow, \\& Junkkarinen 1997). It is also unlikely that PDLAs are generally due to the QSO host galaxy, again because of the significant velocity differences involved. It therefore seems highly plausible that PDLAs are either members of the same family as, or at least strongly resemble, the population of intervening absorbers. If this interpretation holds true, then investigating the hitherto unexploited sample of PDLAs may yield further insights into the connection between AGN activity and galaxy formation. Based on a radio-selected sample of QSOs, CORALS (Ellison et al. 2002) is one of the most recent DLA surveys aimed at determining the extent of any dust bias affecting the statistics of DLAs drawn from optically-selected, magnitude limited QSO samples. In this paper, we present the PDLAs found by the CORALS survey (\\S2) and compare their rate of incidence with that of intervening absorbers at $\\Delta v > 3000$\\,km~s$^{-1}$ in the same sample of QSOs (\\S3). We also compare the frequency of PDLAs between RLQs and RQQs. Our main result is that we find a possible excess of PDLAs towards RLQs (currently only confirmed at slightly below the $2\\sigma$ level), and that since all of the systems that we identify lie in front of optically bright QSOs it appears that PDLAs contain only a modest amount of dust. These findings, which are still tentative because of the small size of present samples, are briefly discussed in \\S4. ", + "conclusions": "The main aim of this paper has been to assess whether PDLAs can provide an insight into the the clustering of galaxies around QSOs at high redshifts. The main conclusions of this work may be summarised as follows. \\begin{enumerate} \\item Despite the small number statistics of this work (4 DLAs in the combined CORALS + FBQS sample), we find evidence that our radio-selected sample of QSO exhibits an excess (by a factor of 4) of PDLAs at \\zae\\ compared to intervening redshifts. This implies an excess of material (presumably galaxies) near the QSO at levels consistent with faint galaxy excesses around RLQs (e.g. Sanchez \\& Gonzalez-Serrano 1999). At present this result is only significant at the $\\sim 2 \\sigma$ level, due to the limited statistics. However, it is certainly suggestive and emphasises the potential use of proximate DLAs as probes of the environment of QSOs. \\item Similarly, the number density of PDLAs towards RLQs is $> 4$ times higher than that of PDLAs towards RQQs. In fact, we find that $n(z)$ of PDLAs towards RQQs is indistinguishable from that of intervening DLAs. This result is consistent with the premise that RLQs preferentially mark cluster environments. \\item Since all of the PDLAs in the CORALS sample are found towards optically bright QSOs ($ B < 20$) there appears to be no strong extinction effect from dust. Although previous studies find that dust in \\zae\\ systems may redden the background QSO (Carilli et al 1998), it appears that the amount of dust that is present is modest and not sufficient to exclude a significant number of quasars from optical surveys. \\item We identify a new PDLA in the line of sight towards B0405$-$331 which may exhibit \\lya\\ emission, based on residual flux seen at the base of the saturated absorption trough. \\end{enumerate} Lacy et al. (2001) have recently proposed a unified scheme for radio-loud and radio-quiet QSOs by showing that the radio luminosity of QSOs scales with the black hole mass and the accretion rate onto the black hole. Although the data support the broad increase in radio luminosity with black hole mass, it is unclear whether accretion rate is the physical driver of this relation (Dunlop et al. 2002). We have presented in this paper evidence suggesting that radio-loud QSOs are found preferentially in richer environments than radio-quiet QSOs. At first sight, such a connection between the pc scale of the central engine and the Mpc scale of galaxy clusters may seem surprising. However, it may simply be another manifestation of the relationship uncovered by Magorrian et al. (1998) between black hole mass and bulge mass, if the most massive galaxies are also more likely to be found in rich clusters. With the current limited statistics for PDLAs it is difficult to extend the analysis of the DLA environment around QSOs. As the number of identified \\zae\\ systems increases (several candidates should be present in existing surveys) it will be possible to compare the 1-D QSO-DLA correlation function with the galaxy-QSO correlation function in clusters (Mart\\'{i}nez et al. 1999). Such a comparison will provide further clues to the nature of PDLAs, and in turn contribute greatly to the investigation of QSO cluster environments. Of course, we must bear in mind that the 3000 \\kms\\ cut-off adopted as the definition for associated DLAs is somewhat arbitrary. This point may be especially important since one of the CORALS PDLAs is right on the 3000 \\kms\\ cut-off. With improved statistics, it will be very interesting to investigate $n(z)$ as a function of relative velocity from the QSO." + }, + "0112/astro-ph0112073_arXiv.txt": { + "abstract": "{We have investigated the nature of the magnetic white dwarf \\lp\\ = LHS\\,2293 by polarimetric monitoring, searching for short-term variability. No periodicity was found and we can exclude rotation periods between 4 sec and 1.5\\,hour with a high confidence. Maximum amplitudes of sinusoidal variations are $\\Delta R < 0.009$\\,mag and $\\Delta V_{\\rm R} < 0.7$\\,\\% for a mean value of the $R-$band circular polarization of $V_{\\rm R} = +9.1\\pm0.3$\\,\\%. Combined with earlier results by other authors, our observation suggests that \\lp{} is, in fact, an extremely slowly rotating single white dwarf and not an unrecognized fast rotator and/or disguised cataclysmic variable. ", + "introduction": "The large majority of white dwarfs are slow rotators with equatorial velocities \\vrot\\,sin\\,i~$<\\,15$\\,km\\,s$^{-1}$ and rotational periods \\prot\\,$ \\ga 1$\\,hr (Heber et al. 1997, Koester et al. 1998). Much slower rotation is not detectable in non-magnetic white dwarfs, but easily measureable in magnetic ones by polarimetric monitoring (Schmidt \\& Norsworthy 1991, Berdyugin \\& Piirola 1999). Among the magnetic white dwarfs, there is a surprising dichotomy in the distribution of \\prot\\ for magnetic white dwarfs with all stars having either \\prot\\,$<\\,20$\\,d or \\prot\\,$>\\,100$\\,yrs (Schmidt \\& Norsworthy 1991). Some magnetic white dwarfs rotate surprisingly fast while others are apparently extremely slow. Among the fast ones are the DA white dwarf RE0317-853 with $B \\simeq 500$ MG, \\teff\\, $\\simeq 40\\,000$\\,K, and \\prot\\,= 12 min (Barstow et al. 1995, Ferrario et al. 1997, Burleigh et al. 1999), the DA star PG1015+014 with $B \\simeq 120$ MG, \\teff\\,$\\simeq 14\\,000$\\,K, and \\prot\\,= 99 min (Wickramasinge \\& Cropper 1988, Schmidt and Norsworthy 1991), and the DAB white dwarf Feige\\,7 with $B \\sim 35$ MG, \\teff\\,$\\simeq 20\\,000$\\,K, \\prot\\,= 2.2\\,hr (Liebert et al. 1977, Achilleos et al. 1992). Five systems seem to be very slow rotators with \\prot\\,$ > 100$\\,yr (Schmidt \\& Norsworthy 1991), among them the proven systems, GD229, G240-72 (Berdyugin \\& Piirola 1999), and Grw+$70^{\\circ}8247$ (Friedrich \\& Jordan 2001), as well as a suspected one, LP790-29 (Liebert et al. 1978). Slow rotation may be caused by coupling of angular momentum into the giant envelope of the progenitor star or the interstellar medium during later stages (Schmidt \\& Norsworthy 1991). Fast rotation may be achieved in a double degenerate which ends as a merger (e.g. Segretain et al. 1997) or in a magnetic cataclysmic variable which loses synchronism (Meyer \\& Meyer-Hofmeister 1999). If the donor in a mass-transfer binary is nearly substellar and hydrogen rich (short-period AM Herculis binary), the white dwarf is expected to be of spectral type DA which is not the case for LP790-29. Cataclysmic variables with a (partially) degenerate low-mass donor (AM CVn binaries) transfer helium or carbon and typically end as CO white dwarfs, possibly with a substellar companion (Iben \\& Tutukov 1991). The hot white dwarf RE0317-853 has been suggested to be the result of a merger (Barstow et al. 1995, Ferrario et al. 1997) or a mass-transfer binary (Meyer \\& Meyer-Hofmeister 1999). The former appears more likely because the primary in RE0317-853 is hot with \\teff\\,$\\simeq 40\\,000$ K and all white dwarfs in short-period cataclysmic variables are cool with \\teff\\,$\\simeq 9\\,000 - 15\\,000$\\,K (G\\\"ansicke 2000). Although definite conclusions in any individual case may be problematic, the detection of rapid rotation of magnetic white dwarfs would clearly help to elucidate their evolutionary history. It also helps to understand the physical processes by which angular momentum is coupled into the environment of the star (Schmidt \\& Norsworthy 1991) and may help to define sources of gravitational wave radiation (Heyl 2000). \\begin{figure*}[t] \\begin{center} \\includegraphics[width=12cm]{h3174.f1} \\caption{Time series of the Stokes intensity $I_{\\rm R}$ and circular polarization $V_{\\rm R}$ in the Bessell $R$-band. The time resolution is \\mbox{$\\sim 30-40$\\,sec} with 150 individual CCD images taken over 1.5\\,h. The $R$-band magnitude was measured relative to two comparison stars in the field of \\lp.} \\label{lightcurves} \\end{center} \\end{figure*} \\lp\\,=\\,LHS\\,2293 was discovered by Liebert et al. (1978) and found to be a highly circularly polarized cool white dwarf which shows the Zeeman shifted C$_2$ Swan bands. It has Stokes $V\\,\\simeq\\,+8$\\% and +9\\% at wavelengths shortward of 4300\\,\\AA\\ and longward of 5500\\AA, respectively, (Liebert et al. 1978, West 1989) which decreases to nil inbetween and to 2\\% in the J--band. Bues (1999) refined its temperature to \\teff \\,$\\simeq 7800$\\,K. The field strength was originally quoted as $B \\sim 200$\\,MG (Liebert et al. 1978, Schmidt \\& Smith 1995), while Bues used 50\\,MG in her spectral fitting, and Wickramasinghe \\& Ferrario (2000) quoted an uncertain 100\\,MG. \\lp\\ is also linearly polarized at the level of $\\sim 1$\\,\\% (West 1989). Polarimetric observations (Liebert et al. 1978, Robert \\& Moffat 1989, West 1989) have shown that the level of circular polarization has stayed constant over 10 years, excluding rotation periods in the range of $\\sim $20\\,min\\,$\\la P_{\\rm rot}\\,\\la\\,100$\\,yrs (Schmidt \\& Norsworthy 1991). In a non-axisymmetric field geometry, the circular polarization will depend on rotation phase and a short rotation period should be readily detectable in time series of the circular polarization, provided the magnetic axis is inclined for more than a few degrees against the rotational axis and the latter does not point directly at the observer. In this communication, we report results of a search for rapid rotation in \\lp, using photometric and polarimetric data taken in the Bessell $R-$band. Given the observed spectral dependence of Stokes $V$ (Liebert et al. 1978, Schmidt et al. 1995), the $R-$band provides the best polarimetric signal of the standard photometric bands. ", + "conclusions": "We have performed a sensitive search for short periodicities in the $R$-band flux and circular polarization of the highly magnetic white dwarf LP790-29 which was previously thought to have a rotational period \\prot\\ $> 100$\\,yrs (Schmidt \\& Norsworthy 1991). We undertook this search because the sparse data available so far may have prevented the discovery of short periods. Rapid rotation would be the signature of a white dwarf spun up in a cataclysmic variable of the AM Herculis type (Meyer \\& Meyer-Hofmeister (1999), of the AM Canes Venaticorum type (Iben \\& Tutukov 1991), or in a double degenerate merger (e.g. Segretain et al. 1997, Ferrario et al. 1997). A merger may rotate at the disruption limit. Segretain et al. argued, however, that the merger loses 90\\% of the initial angular momentum by a strong wind, yielding an initial rotational period \\prot\\,$\\sim 1$\\,min. At an effective temperature of about 8000\\,K (Liebert et al. 1978, Bues 1999) the cooling age of \\lp\\ is $t_{\\rm cool} \\simeq 2 \\times 10^9$\\,yrs (Anselowitz et al. 1999), but may be shorter if it originated from the merger of a cool white dwarf with its companion (Segretain et al. 1997). Hence, even a cool white dwarf might still be a fast rotator, although over the time the original rotational velocity may have been reduced by magnetic braking. Our principal result is the absence of variability in \\lp\\ with periods between 4\\,sec and about 3 hours and amplitudes $\\Delta R > 0.009$\\,mag and $\\Delta V_{\\rm R} > 0.7$\\,\\%. This includes the absence of photometric variability of the type one might expect in a short period binary. Hence, there is no positive evidence for fast rotation and no evidence for any of the above scenarios. The only remaining possibilities which could mask rapid rotation in \\lp\\ are (i) the almost perfect alignment of the rotational and magnetic axes of an \\mbox{azimuthally} symmetric field or (ii) a rotational axis oriented directly towards the observer, leading to rotational variability below our detection limit of the circular polarization. Previous circular polarimetry excludes periods longer than \\mbox{\\prot\\ $\\sim 1$\\,h}, although the limit \\prot\\ $> 100$\\,yrs set by West (1989) and Schmidt \\& Norsworthy (1991) may be premature in view of the low level of the 1994 circular polarization by Schmidt et al. (1995). Nevertheless, these results suggest that \\lp\\ is, in fact, an exceedingly slow rotator. It seems worthwhile to follow up the possibility of a period of about a quarter of a century (see also the paper by Jordan \\& Friedrich 2001) by monitoring the level of circular polarization." + }, + "0112/astro-ph0112559_arXiv.txt": { + "abstract": "The study of gamma-ray burst (GRB) host galaxies in the radio, sub-mm, and X-ray wavelength regimes began only recently, in contrast to optical studies. This is mainly due to the long timescale on which the radio afterglow emission decays, and to the intrinsic faintness of radio emission from star-forming galaxies at $z\\sim 1$, as well as source confusion in sub-mm observations; X-ray observations of GRB hosts have simply not been attempted yet. Despite these difficulties, we have recently made the first detections of radio and sub-mm emission from the host galaxies of GRB\\,980703 and GRB\\,010222, respectively, using the VLA and the SCUBA instrument on JCMT. In both cases we find that the inferred star formation rates ($\\sim 500$ M$_\\odot$) and bolometric luminosities (${\\rm few}\\times 10^{12}$ L$_\\odot$) indicate that these galaxies are possibly analogous to the local population of Ultra-Luminous Infrared Galaxies (ULIRGs) undergoing a starburst. However, there is a modest probability that the observed emission is due to AGN activity rather than star formation, thus requiring observations with Chandra or XMM. The sample of GRB hosts offers a number of unique advantages to the broader question of the evolution of galaxies and star formation from high redshift to the present time since: (i) GRBs trace massive stars, (ii) are detectable to high redshifts, and (iii) have immense dust penetrating power. Therefore, radio/sub-mm/X-ray observations of GRB hosts can potentially provide crucial information both on the nature of the GRB host galaxies, and on the history of star formation. ", + "introduction": " ", + "conclusions": "" + }, + "0112/astro-ph0112067_arXiv.txt": { + "abstract": "We report the detection of a very narrow P Cygni profile on top of the broad emission \\ha\\ and \\hb\\ lines of the Type IIn Supernova 1997eg. A similar feature has been detected in SN~1997ab (Salamanca et al. 1998), SN~1998S (Meikle \\& Geballe 1998, Fassia et al. 2001) and SN~1995G (Filippenko \\& Schlegel 1995). The detection of the narrow P Cygni profile indicates the existence of a dense circumstellar material (CSM) into which the ejecta of the supernova is expanding. From the analysis of the spectra of SN~1997eg we deduce (a) that such CSM is very dense (n\\gapprox 5\\ET{7} \\cc), (b) that has a low expanding velocity of about 160 \\kms. The origin of such dense CSM can be either a very dense progenitor wind (\\.M \\si \\E{-2} \\msunyr) or a circumstellar shell product of the progenitor wind expanding into a high pressure environment. ", + "introduction": "SN~1997eg is a Type IIn supernova discovered on 1997 December 5 (Nakano \\& Masakatsu 1997) having an unfiltered CCD magnitude of 15.6. Its coordinates are R.A. $=$ 13\\hours 11\\minutes 36\\fseconds73 and DEC $=$ +22\\arcdeg 55\\arcmin 29\\farcs4 (equinox 2000.0), which is 4\\farcs1 west and 33\\farcs1 north of the center of the host galaxy NGC~5012, a spiral galaxy with morphological type SAB(rs)c that host a low luminosity AGN in its center, and is situated at 50 Mpc (Ho, Filippenko, Sargent 1997). From the analysis of a spectrum taken 15 days later, Filippenko \\& Barth (1997) report that in the optical range, SN~1997eg show the typical features of Type IIn SN: absence of broad P Cygni profiles and, instead, strong emission lines, notably \\ha\\ and \\hb\\ lines, on top of a very blue continuum - like most Type IIn supernovae. The He{\\smc I} 5876 \\AA\\ is very strong, suggesting either a very high Helium abundance or a blend with the Na{\\smc I} 5894 \\AA\\ blend. From the ratio of the [O{\\smc III}] 4363/5007 lines they deduce the presence of a very dense circumstellar material (n \\gapprox \\E{8} \\cc). Besides, lines of very high excitation like He{\\smc II} 4886 \\AA\\ or [Fe{\\smc X}] 6375 \\AA\\ are prominent (Filippenko \\& Barth 1997) indicating the presence of hard radiation. Finally, SN~1997eg has been detected in radio with the Very Large Array (VLA). The flux measured was 0.52 $\\pm$ 0.06 mJy at 3.6 cm on 1998 May 31 and 0.53 $\\pm$ 0.12 on 1998 June 9 (Lacey \\& Weiler 1998). The exact date of explosion of SN~1997eg is not known, only that it was not seen on 1997 August 11 (\\si 4 months before its discovery). We will adopt the discovery date as the date of explosion. Therefore, our data were taken when the supernova had an age of about 200 days. ", + "conclusions": "There is strong observational evidence for the existence of a dense and hot CSM in Type IIn SN, particularly in SN~1997eg. Such dense environment was ``necessary'' from the theoretical point of view, because the physical explanation invokes the presence of radiative shocks, produced via the interaction of the SN ejecta with dense material. Echelle spectrum of SN~1997eg shows a very narrow P Cygni line atop the broad emission on \\ha\\ and \\hb. This feature seems to be common in Type IIn SN in their early stages, and points to either a massive and slow wind of the progenitor just prior to the explosion or to a wind in a high pressure medium, as its origin. However, even though the data presented here answer many questions concerning such CSM, it opens even more, which can only be addressed with more data. Instead of concentrating in small details of SN IIn, long-term, multi-wavelength monitoring of such objects would be the right thing to do (e.g. Aretxaga et al. 1998) in a similar way as is done for AGN (see for example Peterson \\& Wandel 2000)." + }, + "0112/astro-ph0112251_arXiv.txt": { + "abstract": "In an effort to improve the utility of the helium burning red clump luminosity as a distance indicator, we explore the sensitivity of the K-band red clump absolute magnitude [$M_K(RC)$] to metallicity and age. We rely upon JK photometry for 14 open clusters and two globulars from the 2nd Incremental Data Release of the 2MASS Point Source Catalog. The distances, metallicities, and ages of the open clusters are all on an internally consistent system, while the K(RC) values are measured from the 2MASS data. For clusters younger than $\\sim$2 Gyr, $M_K(RC)$ is insensitive to metallicity but shows a dependence on age. In contrast, for clusters older than $\\sim$2 Gyr, $M_K(RC)$ is influenced primarily by the metallicity of the population and shows little or no dependence on the age. Theoretical red clump models based on the formalism of Girardi et al. reinforce this finding. Over comparable metallicity and age ranges, our average $M_K(RC)$ value is in accord with that based on solar-neighborhood red clump stars with HIPPARCOS parallaxes. Lastly, we compute the distance to the open cluster NGC 2158 using our red clump calibration. Adopting an age of $1.6 \\pm 0.2$ Gyr and $[Fe/H] = -0.24 \\pm 0.06$, our calibration yields a distance of $(m-M)_{V} = 14.38 \\pm 0.09$. ", + "introduction": "During the past few years, the helium burning red clump (RC) has gained considerable attention for its potential as a standard candle. The primary advantage of the RC is the ease with which it can be recognized in the color-magnitude diagram. However, there is currently a great deal of controversy in the literature regarding the appropriate treatment of possible metallicity and age effects on the I-band absolute magnitude of the RC ($M_I(RC)$). There are two schools of thought on this issue; the first assumes a constant value for $M_I(RC)$ which is then used to facilitate a single-step distance determination via knowledge of the apparent RC magnitude and the extinction (e.g. Paczy\\'nski \\& Stanek 1998; Stanek \\& Garnavich 1998). The second approach is founded on the claim that both age and metal abundance have a significant influence on the luminosity of RC stars (e.g. Cole 1998; Sarajedini 1999) and must be accounted for in determining $M_I(RC)$ and therefore the distance. Both Paczy\\'nski \\& Stanek (1998) and Stanek \\& Garnavich (1998) use HIPPARCOS RC stars with parallax errors of less than 10\\% to calculate the I-band absolute magnitude of the solar neighborhood red clump. In their analysis, Paczy\\'nski \\& Stanek (1998) find that $M_{I}(RC)$ shows no variation with color over the range $0.8 < (V-I)_{0} < 1.4$ and, from a Gaussian fit to the RC luminosity function, find $M_{I}(RC) = -0.28 \\pm 0.09$. Following the same methodology and building upon the earlier work, Stanek \\& Garnavich (1998) find a similar result with $M_{I}(RC) = -0.23 \\pm 0.03$. With this calibration, a single step calculation is then used to determine the distance to the Galactic center (Paczy\\'nski \\& Stanek 1998) and M31 (Stanek \\& Garnavich 1998). Both of these investigations found little or no variation in $M_I$ of the RC stars with color; this was taken to imply that $M_{I}(RC)$ does not vary significantly with metallicity. In contrast, theoretical models from Girardi \\& Salaris (2001) and the earlier models of Seidel, Demarque, \\& Weinberg (1987; see also Cole 1998) show that $M_I(RC)$ is dependent on both age and metallicity, becoming fainter as both increase. These models are in good agreement with the observations presented by Sarajedini (1999, hereafter S99). Using published photometry for 8 open clusters, S99's most important result is that while $M_{I}(RC)$ is less sensitive to metal abundance than $M_{V}(RC)$, both still retain a considerable dependence on the age and metallicity of the stellar population. As a result, the single-step method of applying the solar-neighborhood $M_{I}(RC)$ to populations with a different age-metallicity mix could be problematic. Alves (2000) also uses the HIPPARCOS RC for his calibration; however, he relies upon the K-band luminosity ($M_K$) of the RC stars in the hope that, since the K-band is less sensitive to extinction (and possibly metallicity as well) than the I-band, it might make a better choice as a standard candle. Alves (2000) restricts his RC stars to those that have metallicities from high resolution spectroscopic data. For this group of 238 RC stars, he finds a peak value of $M_K(RC) = -1.61 \\pm 0.03$ with no correlation between $[Fe/H]$ and $M_K$. However, he is not able to explore the effect of age on $M_K(RC)$ due to the lack of such information for the individual stars in his sample. These previous works have prompted us to combine the approaches of S99 and Alves (2000) and to investigate the influence of age and metal abundance on $M_K(RC)$ for a number of open clusters with well-known distances and metallicities. In section 2 we discuss the observational data. Section 3 compares our data with the results of theoretical models and presents a discussion of the results; our conclusions are summarized in Section 4. ", + "conclusions": "In this paper, we have sought to establish the K-band absolute magnitude of the helium burning red clump stars ($M_K(RC)$) as a distance indicator. To facilitate this, we have utilized infrared photometry from the 2MASS catalog along with distances, metallicities, and ages for 14 open clusters and 2 globular clusters. Our sample encompasses an age range from 0.63 Gyr to 12 Gyr and metallicities from --1.15 to 0.15 dex. Based on an analysis of these data, we draw the following conclusions. 1. There is a statistically significant range of $M_K(RC)$ values among the star clusters in our sample. In particular, for the 14 open clusters, we calculate $\\langle$$M_{K}(RC)$$\\rangle$$= -1.62$ with a standard deviation of 0.21 mag. In contrast, the mean error in these $M_K(RC)$ values is 0.13 mag. 2. Upon inspection of figures 5 and 6, we find that for clusters younger than $\\sim$2 Gyr, $M_K(RC)$ is insensitive to metallicity but shows a dependence on age. In contrast, for clusters older than $\\sim$2 Gyr, $M_K(RC)$ is influenced primarily by the metallicity of the population and shows little or no dependence on the age. 3. In general, $M_K(RC)$ is less sensitive to age and metallicity than $M_I(RC)$ over the parameter range common to both this paper and Sarajedini (1999) from which the $M_I(RC)$ values are taken. 4. Over comparable metallicity and age ranges, our average $M_K(RC)$ value of --1.62 mag is consistent with that of Alves (2000) which is based on solar-neighborhood red clump stars with HIPPARCOS parallaxes. We also suggest that the significant scatter in the Alves (2000) $M_K$ data is likely due to a range of ages between $\\sim$1.6 and $\\sim$4 Gyr among these stars. 5. The theoretical red clump models based on the formalism of Girardi et al. (2000) agree reasonably well with our observational data, indicating that age plays an important role in determining $M_K(RC)$ for younger populations while metallicity mainly affects older populations. 6. Using the K-band absolute magnitude of the red clump, we are able to compute the distance to the open cluster NGC 2158. Adopting an age of $1.6 \\pm 0.2$ Gyr and $[Fe/H] = -0.24 \\pm 0.06$, our calibration yields a distance of $(m-M)_{V} = 14.38 \\pm 0.09$. 7. When determining distances for star clusters having $-0.5 \\leq [Fe/H] \\leq 0.0$ and $10^{9.2} \\leq age \\leq 10^{9.9}$, one can ignore the interpolation discussed in section 3.3 and simply use $$ = $-1.61 \\pm 0.04$." + }, + "0112/astro-ph0112298_arXiv.txt": { + "abstract": "A correlation is presented between the pulse lag and the jet-break time for seven BATSE gamma-ray bursts with known redshifts. This is, to our best knowledge, the first known direct tight correlation between a property of the gamma-ray burst phase (the pulse lag) and the afterglow phase (the jet-break time). As pulse lag and luminosity have been found to be correlated this also represents a correlation between peak luminosity and jet-break time. Observed timescales (variability or spectral lags) as well as peak luminosity naturally have a strong dependence on the Lorentz factor of the outflow and so we propose that much of the variety among GRBs has a purely kinematic origin (the speed or direction of the outflow). We explore a model in which the variation among GRBs is due to a variation in jet-opening angles, and find that the narrowest jets have the fastest outflows. We also explore models in which the jets have similar morphology and size, and the variation among bursts is caused by variation in viewing angle and/or due to a velocity profile. The relations between luminosity, variability, spectral lag and jet-break time can be qualitatively understood from models in which the Lorentz factor decreases as a function of angle from the jet axis. One expects to see high luminosities, short pulse lags and high variability as well as an early jet-break time for bursts viewed on axis, while higher viewing inclinations will yield lower luminosities, longer pulse lags, smoother bursts and later jet-break times. ", + "introduction": "Pulse peaks in Gamma-Ray Burst (GRB) lightcurves evolve in time from higher to lower energies and become wider. \\citet{nmb00} provide measurements of such lags as observed in GRBs between different energy channels of the Burst And Transient Source Experiment (BATSE). The lags are measured by calculating the cross-correlation function (CCF). In a sample of GRBs for which the redshift is known it is found that the spectral lag is related to the isotropic gamma-ray peak luminosity \\citep{nmb00}. Similarly \\citet{fr00} have shown that measures of variability in GRB lightcurves (see also Reichart et al. [2001]\\nocite{rlfr+01}) are related to the isotropic gamma-ray peak luminosity. Vice versa, measurements of variability or of spectral lags can then be used as a crude distance indicator for the GRB. Observed timescales (variability or spectral lags) as well as peak luminosity naturally have a strong dependence on the Doppler factor of the outflow (a function of the Lorentz factor and the direction of motion with respect to the observer). If indeed the Doppler factor is the dominant parameter among GRBs, a relation between spectral lags/variability and luminosity is expected. Crucial to our understanding of what causes the GRB is the question of whether GRB engines are in some sense ``standard candles''. Observationally it is found that the isotropic equivalent energies of GRBs range from about 5 $\\times$ 10$^{51}$ to 1.4 $\\times$ 10$^{54}$ ergs \\citep{bfs01}. However, transitions have been observed at optical and radio wavelengths which can be interpreted as being due to collimated (jetted) outflow \\citep{hbf+99}. When correcting the observed $\\gamma$-ray energies for the geometry of the outflow, GRB energies appear narrowly clustered around $5 \\times 10^{50}$ ergs (Frail et al. 2001; see also Panaitescu and Kumar 2001\\nocite{fksd01,pk01}. \\citet{fksd01} infer the jet opening angle from the observed jet-break time and find that there is a wide range in opening angles. The reason for why this range in angles exists is currently not understood. The idea that GRB energies may be narrowly clustered was also put forward by \\citet{jay00,jay01}. However, \\citet{jay01} proposed that there exists not a range in opening angles, but that all bursts derive from a single-burst-jet morphology and that the variation in viewing angle of the jet yields the observed variation among GRBs. In this paper is presented a tight correlation between jet-break times, $\\tau_j$, and pulse lags, $\\Delta t$. Since luminosity and pulse lags have been shown to be correlated, this also represents a relation between luminosity and jet-break times. As time scales and peak luminosity are strong functions of the Doppler factor of the outflow, this relation suggests the jet-break time to also be a strong function of the Doppler factor. We discuss the implications of this result on our understanding of the morphology of the explosion. ", + "conclusions": "Herein we have presented a tight relation between GRB pulse lags and afterglow jet-break times (Eqn.~\\ref{t_jeqn}). As spectral lag was shown to be related to the isotropic gamma-ray peak luminosity \\citep{nmb00} this also represents a correlation between peak luminosity and jet-break time. This is, to our best knowledge, the first known direct correlation between a property of the gamma-ray burst phase (the pulse lag or peak luminosity) and the afterglow phase (the jet-break time). As observed timescales (variability or spectral lags) as well as peak luminosity naturally have a strong dependence on the Lorentz factor, $\\gamma$, or angle with respect to the motion of the outflow, $\\theta$, we propose that the variety among GRBs has a purely kinematic origin. The emergence of simple trends between jet-break time, pulse lag and luminosity gives us clues and constraints on how such a model can be constructed. \\citet{fksd01} infer that the differences among observed GRB energies are due to a range of jet opening angles. Within this framework we have then shown that the fastest GRBs, with the highest $\\gamma$, have the narrowest jets (Sec.~\\ref{burstpop}). A prediction of this model is that not only should there be a variation in opening angles $\\theta_j$ among bursts, but they should be related to Lorentz factor $\\gamma$ as $\\theta_j \\propto \\gamma^{-3/8}$, and jet angles $\\theta_j$ have to be produced with a probability $P \\sim P_{obs}/\\theta_j^2 \\propto \\theta_j^{-4.54}$. We then have explored the possibility that variation among GRBs, all now assumed to be morphologically the same, is caused by variation in observer viewing angle from the jet axis (Sec.~\\ref{simplejet}), or Lorentz factor of the jet as a function of angle from the jet axis (Sec.~\\ref{structuredjet}). These models qualitatively show that variation of the observed Doppler factor, ${\\cal D}$, (Eqn. \\ref{deltaeqn}) from a single jet morphology can produce the observed variations among GRB luminosities and timescales: $L_{pk} \\propto N_{pk} \\propto 1/\\Delta t \\propto 1/\\tau_j \\propto {\\cal D}$. These models also provide a natural explanation for the probability distribution of observation angles $P_{obs}$ consistent with that reported by \\citet{fksd01}. A realistic model of a single-burst morphology will have both perspective effects as in Section \\ref{simplejet} and jet-structure effects as in Section \\ref{structuredjet}. Set inside the collapsar progenitor model \\cite[e.g.][]{mcfw99} a cohesive picture begins to emerge. A jet driven out from the center of the star will vary as a function of angle from the jet axis. Ejecta at the center will be faster and lighter and, bearing the brunt of the as yet unknown acceleration mechanism, will be more fractured. At larger angles from the jet core, ejecta will interact with the stellar wall, and thus will be slower and will entrain more baryons. Thus one expects to see high luminosities, short pulse lags and high variability as well as an early jet break time for bursts viewed on axis, while higher viewing inclinations will yield lower luminosities, longer pulse lags, smoother bursts and later jet break times. Thus the variability-luminosity relationship \\citep{rlfr+01} as well as the spectral lag-jet break time and the spectral lag-luminosity (Eqns.~\\ref{t_jeqn} \\& \\ref{jpnlageqn}) can be naturally accomodated." + }, + "0112/astro-ph0112121_arXiv.txt": { + "abstract": "Two lines of thought exist as to the nature of Soft Gamma-ray Repeaters (SGRs) and Anomalous X-ray Pulsars (AXPs). On the one hand, Duncan \\& Thompson (1992) and Thompson \\& Duncan (1995) propose neutron stars with super-critical ($>10^{14}$ G) magnetic fields, which spin-down the stars and power the gamma-ray bursts. On the other hand, several authors (van Paradijs, Taam \\& van den Heuvel 1995; Chatterjee, Hernquist \\& Narayan 2000; Alpar 2001; Marsden et al. 2001; Menou, Perna \\& Hernquist 2001) propose neutron stars with typical pulsar magnetic fields ($\\sim10^{12}$ G), which are spun-down by magnetospheric ``propeller'' torques from fallback or fossil disks in addition to magnetic dipole radiation. We discuss these two concepts in light of various observations. ", + "introduction": "Magnetars, defined to be neutron stars possessing dipole magnetic fields in excess of the quantum critical value of 4.4$\\times$10$^{13}$ Gauss, constitute a proposed class distinct from radio and x-ray pulsars, in which magnetic energy, rather than rotational energy, plays the dominant role in powering emissions. The strong magnetic dipole radiation (MDR) would spin-down magnetars quite rapidly leaving them with spin periods of a few seconds after $\\sim$10$^3$ years. Repetitive soft gamma-ray bursts are interpreted as due to crust cracking events in the neutron star surface, whereas the super busts seen from SGR0525-66 and SGR1900+14 would result from sudden large-scale magnetic reconnection. Problems replicating the estimated ages of SGR/AXPs in this model have led to modeling extra sources of torque on the system, but with the magnetic energy remaining as the dominant power source (Kouveliotou et al. 1999). Alternatively, the rapid spindown rates, young ages inferred from the SNR ages, long spin periods clustered around 5-10 s, and $\\sim$ 10$^{35}$ erg/s x-ray luminosities for SGR/AXPs can all be explained by models involving the propeller effect on inflowing material as the dominant spindown torque. This material comes from a small accretion disk formed around the neutron star very early in its life. Such a disk can form in several ways: from the inner most ejecta material falling back within a few hours of the initial supernova explosion (Michel 1988; Chatterjee et al 2000); from the reversal of slower-moving inner ejecta by the Sedov phase reverse shock relatively soon after the blast wave hits the progenitor winds (Truelove \\& McKee 1999); or from high velocity neutron stars capturing comoving ejecta (van Paradijs et al. 1995). Only a very small fraction of the ejecta is needed to form a fossil disk of 10$^{-6} M_\\odot$ which is all that is required to explain the spindown of SGR/AXPs via the propeller mechanism. In this model, the exceedingly rapid spindown causes crust cracking and subduction to provide both the energy and mechanism for the very energetic bursts. ", + "conclusions": "The success of the accretion models is that they require only the well-studied properties of neutron stars and supernovae, and they can be applied beyond AXP/SGRs to clarify contradictions in interpretations of other neutron stars. These models predict the non-bursting attributes --- luminosity, spin period, spin-down rate --- as well as the low number seen in the Galaxy. Spin-down driven quakes can also power both the repetitive bursting and the super bursts, and the durations of these bursts are consistent with postquake vibrational damping times. Direct observations of the disks, however, are needed to establish their existence. The magnetar model with relativistic winds can also explain both the persistent and bursting x- and gamma-ray emission from SGRs, and the spin-down of both the SGRs and AXPs, if the wind x-ray emission efficiency is near 100\\%. The magnetar model, however, does not explain the clustering of spin periods observed in these sources, even with magnetic field decay. Theoretical arguments suggest that magnetic fields can exist far above the critical field, but observational evidence from all of the radio pulsars, whose implied fields from $P$ and $\\dot{P}$ span over 5 orders of magnitude, show a clear cutoff just short of the critical field." + }, + "0112/physics0112093_arXiv.txt": { + "abstract": "We performed accurate calculation of $\\alpha$-dependence ($\\alpha=e^2/hc$) of the transition frequencies for ions, which are used in a search for the variation of the fine structure constant $\\alpha$ in space-time. We use Dirac-Hartree-Fock method as a zero approximation and then the many-body perturbation theory and configuration interaction methods to improve the results. An important problem of level pseudocrossing (as functions of $\\alpha$) is considered. Near the crossing point the derivative of frequencies over $\\alpha$ varies strongly (including change of the sign). This makes it very sensitive to the position of the crossing point. We proposed a semiempirical solution of the problem which allows to obtain accurate results. ", + "introduction": "Recently there was an intensive discussion of the possible space-time variation of the fine structure constant $\\alpha=e^2/hc$ at the cosmological scale. The first evidence for such variation has been reported in \\cite{WFC99,WMF01,MWF01a,MWF01,MWF01c,MWF01d} from the analysis of the astrophysical data. These results are to be compared with the number of experimental upper bounds on this variation obtained from other astrophysical observations (see, e.g. \\cite{CW99,IPV99,VIP99}) and from the precision laboratory measurements \\cite{PTM99,SBN99,SBN00}. Recently a number of new laboratory tests have been proposed (see, e.g. \\cite{BPM01}). The analysis of the microwave background radiation can also give some restrictions on time variation of $\\alpha$ as suggested in \\cite{KST99,Han99,KS00}. Implementations of the space-time variation of the fine structure constant to the theory of the fundamental interactions are discussed e.g. in Refs.~\\cite{AM97,BIR99,Kuh99,Fuj00,CLV01,KR01,SBM01} (see also disscussion and references in \\cite{MWF01a}). The most straitforward way to look for the variation of $\\alpha$ is to measure the ratio of some fine structure interval to an optical transition frequency, such as $\\omega( np_{1/2} \\rtw np_{3/2})$ and $\\omega( n's_{1/2} \\rtw np_{3/2})$~\\endnote{In fact, the frequency $\\omega( np_{1/2} \\rtw np_{3/2})$ is not measured directly, but is found as a difference: $\\omega(n's_{1/2}\\rtw np_{3/2})-\\omega( n's_{1/2} \\rtw np_{1/2})$.}. This ratio can be roughly estimated as $0.2\\,\\alpha^2 Z^2$, where $Z$ is the nuclear charge \\cite{Sob79}. Therefore, any difference in this ratio for a laboratory experiment and a measurement for some distant astrophysical object can be easily converted into the space-time variation of $\\alpha$. However, as it was pointed out in \\cite{DFW99b}, one can gain about an order of magnitude in the sensitivity to the $\\alpha$-variation by comparing optical transitions for different atoms. In this case the frequency of each transition can be expanded in a series in $\\alpha^2$: \\begin{subequations} \\label{i1} \\begin{eqnarray} \\label{i1a} \\omega_i &=& \\omega_i^{(0)} + \\omega_i^{(2)} \\alpha^2 + \\dots \\\\ \\label{i1b} &=& \\omega_{i,\\textrm{lab}} + q_i x + \\dots, \\quad x \\equiv \\left(\\alpha/\\alpha_0\\right)^2-1, \\end{eqnarray} \\end{subequations} where $\\alpha_0$ stands for the laboratory value of the fine structure constant. Note, that \\Eref{i1a} corresponds to the expansion at $\\alpha=0$, while \\Eref{i1b} ---~to the expansion at $\\alpha=\\alpha_0$. In both cases parameters $\\omega^{(2)}_i$ and $q_i$ appear due to relativistic corrections. For a fine structure transition the first coefficient on the right hand side of \\eref{i1a} turns to zero, while for the optical transitions it does not. Thus, for the case of a fine structure and an optical transition one can write: \\begin{eqnarray} \\label{i2} \\frac{\\omega_{\\textrm{fs}}}{\\omega_{\\textrm{op}}} = \\frac{\\omega_{\\textrm{fs}}^{(2)}}{\\omega_{\\textrm{op}}^{(0)}} \\, \\alpha^2 + O(\\alpha^4), \\end{eqnarray} while for two optical transitions $i$ and $k$ the ratio is: \\begin{eqnarray} \\label{i3} \\frac{\\omega_i}{\\omega_k} = \\frac{\\omega_i^{(0)}}{\\omega_k^{(0)}} +\\left(\\frac{\\omega_i^{(2)}-\\omega_k^{(2)}}{\\omega_k^{(0)}} \\right) \\, \\alpha^2 + O(\\alpha^4). \\end{eqnarray} Quite often the coefficients $\\omega_i^{(2)}$ for optical transitions are about an order of magnitude larger than corresponding coefficients for the fine structure transitions $\\omega_{\\textrm{fs}}^{(2)}$ (this is because the relativistic correction to a ground state electron energy is substantially larger than the spin-orbit splitting in an excited state \\cite{DFW99a,DFW99b}). Therefore, the ratio \\eref{i3} is, in general, more sensitive to the variation of $\\alpha$ than the ratio \\eref{i2}. It is also important that the signs of coefficients $\\omega_i^{(2)}$ in \\eref{i3} can vary. For example, for $s$-$p$ transitions the relativistic corrections are positive while for $d$-$p$ transitions they are negative. This allows to suppress possible systematic errors which ``do not know'' about the signs and magnitude of the relativistic corrections \\cite{DFW99b}. On the other hand, for many cases of interest, the underlying atomic theory is much more complicated for \\Eref{i3}. In particular, the most difficult case corresponds to transitions to highly excited states of a multi-electron atom, where the spectrum is very dense. And this happens to be a typical situation for astrophysical spectra, in particular, for large cosmological red shifts. Corresponding atomic calculations have to account very accurately for the electronic correlations, which may affect such spectra quite dramatically. The first calculations of the coefficients $q$ from \\Eref{i1} for the transitions suitable for astronomical and laboratory measurements were done in Refs.~\\cite{DFW99a,DFW99b,DF00a,DFM01}. Here we present a new and more accurate calculations of the coefficients $q$ for the transitions, which are currently used in the analysis of the astrophysical data. A full list of these transitions was given in \\cite{MWF01a}. We have not recalculated here the lightest and the most simple atoms Mg and Al, for which the previous calculation \\cite{DFW99b} should be sufficiently accurate and focused on more complicated ions Si~II, Cr~II, Fe~II, Ni~II, and Zn~II. Our final results for them are given in \\tref{tab_fin}. Note, that here we use the single parameter $q$ instead of two parameters $q_1$ and $q_2$ used in the earlier works and $q \\equiv \\partial \\omega/\\partial x|_{x=0} = q_1+2q_2$. Details of the calculations and discussion of the accuracy will be given in Sec.~\\ref{details}. Before that we briefly address few theoretical points in Sec.~\\ref{theory}. \\begin{table}[tb] \\caption{Final results for parameters $q$ from \\Eref{i1} for Si~II, Cr~II, Fe~II, Ni~II, and Zn~II. Estimated errors are in brackets.} \\label{tab_fin} \\begin{tabular}{llcldrl} \\hline \\hline \\multicolumn{1}{c}{Ion} &\\multicolumn{3}{c}{Transition} &\\multicolumn{1}{c}{$\\omega_0$ (\\cm)} &\\multicolumn{2}{c}{$q$ (\\cm)}\\\\ \\hline Si II &$^2\\!P^o_{1/2}$&$\\rtw$&$^2\\!D_{3/2} $& 55309.3365 &$ 520 $& (30)\\\\ & &$\\rtw$&$^2\\!S_{1/2} $& 65500.4492 &$ 50 $& (30)\\\\ \\hline Cr II &$^6\\!S_{5/2}$& $\\rtw$&$^6\\!P^o_{3/2}$& 48398.868 &$-1360 $& (150)\\\\ & & $\\rtw$&$^6\\!P^o_{5/2}$& 48491.053 &$-1280 $& (150)\\\\ & & $\\rtw$&$^6\\!P^o_{7/2}$& 48632.055 &$-1110 $& (150)\\\\ \\hline Fe II &$^6\\!D_{9/2}$& $\\rtw$&$^6\\!D^o_{9/2}$& 38458.9871 &$ 1330 $& (150)\\\\ & & $\\rtw$&$^6\\!D^o_{7/2}$& 38660.0494 &$ 1490 $& (150)\\\\ & & $\\rtw$&$^6\\!F^o_{11/2}$&41968.0642 &$ 1460 $& (150)\\\\ & & $\\rtw$&$^6\\!F^o_{9/2}$& 42114.8329 &$ 1590 $& (150)\\\\ & & $\\rtw$&$^6\\!P^o_{7/2}$& 42658.2404 &$ 1210 $& (150)\\\\ & & $\\rtw$&$^4\\!F^o_{7/2}$& 62065.528 &$ 1100 $& (300)\\\\ & & $\\rtw$&$^6\\!P^o_{7/2}$& 62171.625 &$-1300 $& (300)\\\\ \\hline Ni II & $^2\\!D_{5/2}$& $\\rtw$&$^2\\!F^o_{7/2}$& 57080.373 &$ -700 $& (250) \\\\ & & $\\rtw$&$^2\\!D^o_{5/2}$& 57420.013 &$-1400 $& (250) \\\\ & & $\\rtw$&$^2\\!F^o_{5/2}$& 58493.071 &$ -20 $& (250) \\\\ \\hline Zn II & $^2\\!S_{1/2}$& $\\rtw$&$^2\\!P^o_{1/2}$& 48481.077 &$ 1584 $& (25) \\\\ & & $\\rtw$&$^2\\!P^o_{3/2}$& 49355.002 &$ 2490 $& (25) \\\\ \\hline \\hline \\end{tabular} \\end{table} ", + "conclusions": "" + }, + "0112/astro-ph0112317_arXiv.txt": { + "abstract": "The formation rate of a close binary consisting of a super-massive black hole and a compact object (presumably a white dwarf) in galactic cusps is calculated with help of the so-called loss cone approximation. For a power-law cusp of radius $r_{a}$, the black hole mass $M\\sim 10^{6} M_{\\odot}$, and the fraction of the compact objects $\\delta \\sim 0.1$ this rate $\\dot N_{wd} \\sim 4\\cdot 10^{-5}K(p)\\sqrt{{GM\\over r_{a}^{3}}} \\approx 3\\cdot 10^{-9}K(p){({M\\over 10^{6}M_{\\odot}})}^{1/2}{({r_{a}\\over 1pc})}^{-3/2}yr^{-1}$. The function $K(p)$ depends on parameter $p$ determining the cusp profile, and for the standard cusp profiles with $p=1/4$ $K(p)\\sim 2$. We estimate the probability ${\\it Pr}$ of finding of a compact object orbiting around a black hole with the period $P$ in one particular galaxy to be ${\\it Pr}\\sim 10^{-7}{({P/10^{3}s\\over M/10^{6}M_{\\odot}})}^{8/3}{({M/ 10^{6}M_{\\odot}\\over r_{a}/ 1pc})}^{3/2}$. The object with the period $P\\sim 10^{3}s$ emits gravitational waves with amplitude sufficient to be detected by LISA type gravitational wave antenna from the distance $\\sim 10^{3} Mpc$. Based on estimates of masses of super-massive black holes in nearby galaxies, we speculate that LISA would detect several such events during its mission. ", + "introduction": "A compact object orbiting around a super-massive black hole with mass $\\sim 10^{5}-10^{6}M_{\\odot}$ with periods $\\sim 10^{3}-10^{4}s$ produces gravitational radiation which could be detected by future space-based gravitational antennas. As was mentioned by a number of authors, such object could settle in a tight orbit around the black hole due to combined action of two body gravitational encounters with other stars in galactic centers and emission of gravitational radiation. The gravitational radiation coming from such object yields a direct information about relativistic field of the black hole. Therefore, it is very interesting to estimate the rate of production of close binaries consisting of a super-massive black hole and a compact object (the capture rate below), and the probability ${\\it Pr}$ of finding of a galaxy with such a binary in the center. This problem has been investigated before by Hils and Bender 1995 and Sigurdsson and Rees 1997. Hils and Bender performed numerical calculation of the capture rate for a black hole and a central stellar cluster with parameters similar to that was observed in the center of M32. Sigurdsson and Rees (hereafter SR) generalized this result by making estimates of the capture rate for a more general model of the galactic centers. The purpose of this paper is to extend the results obtained before by calculation of the capture rate and the probability ${\\it Pr}$ in frameworks of the so-called ``loss cone approximation'' (e. g. Lightman $\\&$ Shapiro 1977, hereafter LS, and references therein). Also, we correct a mistake made in previous estimates. In what follows we consider a black hole and a stellar cluster around it with parameters close to the parameters of the black hole and the stellar cluster in the center of our own Galaxy. At first, we assume that the black hole mass is about $10^{6}M_{\\odot}$. Inside a radius where the black hole mass is approximately equal to the total mass of the stars, a power law increase of the star's number density with decreasing of distance (called the ``cusp'' in the stellar distribution) is predicted by theory (e.g. Bahcall and Wolf 1976). Existence of the cusp is in agreement with observations of the center of our Galaxy (e.g Alexander 1999). Based on observations of our Galaxy and other nearby galaxies, we would expect the cusp radius of the order $\\sim 1pc$. As it will be clear from our discussion (see the next Section), in order to calculate the capture rate we will be interested in inner regions of the cusps, where the density of normal stars is strongly suppressed by star-star collisions. Therefore, we will be interested in compact objects, which could survive collisions with normal stars. There is also another argument for paying special attention to the compact objects in our problem. Namely, a normal star orbiting around a black hole with period of interest must be tidally disrupted, and therefore only the compact objects with small tidal radii may give a persistent source of gravitational radiation. As it was argued by SR, it it reasonable to suppose that the galactic centers contain evolved population of stars, with a large number of white dwarfs (the ratio of the total number of white dwarfs to the total number of the normal stars $\\delta \\sim 0.1$). Since the capture rate is proportional to the total number of compact objects, the white dwarfs could yield the main contribution to the capture rate. Therefore, for simplicity, we consider below the innermost part of the cusp consisting purely of white dwarfs. The generalization of our results on the case of neutron stars and solar mass black holes is straightforward. We calculate the capture rate in Section 2. In Section 3 we estimate the probability ${\\it Pr}$. We summarize and discuss our results in Section 4. G is the Newton constant of gravity and c is the speed of light throughout the Paper. We use expressions from Gradshtein and Ryzhik 1994 when operating with special functions without explicit referencing to that book. ", + "conclusions": "We calculate the formation rate of a close binary in galactic cusps (the capture rate) consisting of a super-massive black hole and a compact object (presumably a white dwarf). We also calculate the probability ${\\it Pr}$ to find a galaxy with such a binary emitting gravitational radiation at frequency $f\\sim 10^{-3}Hz$. We take into account two main process determining the capture rate: 1) two-body gravitational encounters of the compact object, 2) energy loss due to emission of gravitational radiation, and calculate the capture rate analytically, in frameworks of the so-called loss cone approximation. Our main results are given by equations (37),(69) and (80). These equations explicitly show the dependence of the capture rate and probability ${\\it Pr}$ on main parameters of galactic cusps: their radii, the relative number of compact objects in them, the ``sharpness'' $p$ of increase of the number density of compact objects toward the center of a galaxy and the frequency band of gravitational wave antenna receiving the signal. They could be easily used for estimates of the capture rate and probability ${\\it Pr}$ for a broad range of the main parameters. We correct a mistake made by previous researchers in estimate of the capture rate. Our results are obtained in assumption of power law cusps with density $n \\sim r^{-(3/2+p)}$. However, generalization of our results on more realistic cusp profiles (say, with different powers $p$ in the outer and inner regions of the cusp, e.g Alexander 1999) is straightforward. Unfortunately, not so much is known about galactic cusps from observations. Therefore, our results should be considered as qualitative only. Also, much more robust estimates could be obtained using modern computational methods. However, our simple analytical approach clearly shows the dependence on parameters, and therefore, could be used as a guidance for choice of parameters for numerical studies. In their recent review of black holes found in the galactic centers, Kormendy and Gebhardt 2001 identify three black hole with masses of order of $\\sim 10^{6} M_{\\odot}$ (one of them is in our own Galaxy) with distance smaller than $13,2Mpc$. Therefore, it seems reasonable to estimate the density of potential sources of gravitational radiation as $\\ge 10^{-3}Mpc^{-3}$ \\footnote{ This density could be larger due to selection effects. On the other hand, the probability ${\\it Pr}$ is sensitive to the size $r_{a}$ of the cusp (eq. (80)). If all potential sources have the cusps with sizes $\\gg 1pc$, the probability would be suppressed.}. Assuming that we need $\\sim 10^{7}$ potential sources to observe at least one event (equation (80)), we need the sensitivity of gravitational wave antenna to be sufficient to detect a source at the distance $D \\ge 10^{3}Mpc$. The dimensionless amplitude of a gravitational wave coming from a white dwarf orbiting around a black hole with period $\\sim 10^{3}s$ at the distance $\\sim 10^{3}Mpc$ from an observer is of order of $10^{-23}$. This amplitude could be easily detected by a gravitational wave antenna of LISA type. Therefore, we could optimistically expect several events during the life-time of LISA mission $\\sim 3 yr$. Relativistic effects (e.g. influence of Einstein precession and Lense-Thirring precession for a rotating black hole on the shape of the signal) may help to distinguish these events from other sources of gravitational radiation." + }, + "0112/astro-ph0112192_arXiv.txt": { + "abstract": "We analyze the Time-Tagged Event (TTE) data from observations of gamma ray bursts (GRBs) and soft gamma repeaters (SGRs) by the Burst and Transient Source Experiment (BATSE). These data provide the best available time resolution for GRBs and SGRs. We have performed an extensive search for weak periodic signals in the frequency range 400 Hz to 2500 Hz using the burst records for 2203 GRBs and 152 SGR flares. The study employs the Rayleigh power as a test statistic to evaluate the evidence for periodic emissions. We find no evidence of periodic emissions from these events at these frequencies. In all but a very few cases the maximum power values obtained are consistent with what would be expected by chance from a non-periodic signal. In those few instances where there is marginal evidence for periodicity there are problems with the data that cast doubt on the reality of the signal. For classical GRBs, the largest Rayleigh power occurs in bursts whose TTE data appear to be corrupted. For SGRs, our largest Rayleigh power, with a significance of $\\approx 1$\\%, occurs in one record for SGR~1900$+$14 (at $\\approx 2497$~Hz), and in no other outbursts associated with this source; we thus consider it unlikely to represent detection of a real periodicity. From simulations, we deduce that the Rayleigh test would have detected significant oscillations with relative amplitude $\\approx 10$\\% about half the time. Thus, we conclude that high frequency oscillations, if present, must have small relative amplitudes. ", + "introduction": "\\label{sec:motivation} In the last decade observations of classical gamma ray bursts (GRBs) and soft gamma repeaters (SGRs) have finally identified the sites of the sources of these high energy transients, yet the physical nature of the sources remains mysterious (Paczy\\'nski 1995; Costa et al. 1997; Bond 1997; Djorgovski et al. 1997; M\\'esz\\'aros \\& Rees 1997; Horth et al. 1998; Frail et al. 1998; Taylor et al. 1998; Kulkarni et al. 1999). Detailed timing analysis of the high energy emission from these sources offers important clues regarding the nature of the central engines. The positive detection of periodic emission in SGR flares suggests a rotating neutron star origin for these events, but to date, only two events have shown evidence of periodic emission (Barat et al.\\ 1979; Cline et al.\\ 1980; Terrell et al.\\ 1980; Barat et al.\\ 1983; Kouveliotou et al.\\ 1999). The current picture of SGRs was first introduced in 1992 when it was proposed that they are neutron stars with extremely large magnetic fields, coined ``magnetars'' (Duncan and Thompson 1992). Part of the motivation for this model is the 8 s periodicity observed in the 5 March 1979 event, thought to be the rotational period of a neutron star associated with the source. In addition to rotation and precession, faster nonradial pulsations may be involved in the SGR events (McDermott, Van Horn and Hansen 1988; Duncan 1998). To date, no periodic emission has been detected from (classical) GRBs, but it might be expected at some level in some models. If binary compact stars or matter orbiting near neutron stars or black holes are involved in GRBs, it is likely that there are rapid, possibly periodic, associated phenomena. Quasi-periodic oscillations (QPOs) have been observed from accreting neutron star systems with frequencies as high as 1200 Hz. Some QPOs are thought to be associated with gas orbiting a neutron star at the innermost stable circular orbit at $r= 6GM/c^2$, where the Kepler frequency is $2200 (M_\\odot/M)$~Hz for a star of mass $M$; for a black hole it may be as high as $11300 (M_\\odot/M)$~Hz. Recently, significant variability on time scales even shorter than 1~ms has been seen from Cen X-3, attributed to photon bubble oscillations that could produce quasiperiodic variability (Klein et al. 1996; Jernigan, Klein \\& Arons 2000). While it is presently unclear whether there are multiple classes of GRB progenitors, most plausible models involve a central black hole and a (temporary) debris torus around it. This includes massive progenitor systems, such as hypernovae or collapsars where a massive rotating star collapses to a black hole leaving behind an accreting disk of material that releases energy (Paczy\\'nski 1999; Nomoto et al. 2000), and many models involving the merger of compact stars. Short time scale variability, possibly including high frequency ($f\\approx 10^{3}$ Hz) pulsations is expected for these types of sources, where material may be orbiting close to a nascent neutron star or black hole. Further observation of periodic emission in SGR flares or detection of periodic emission from GRBs could provide valuable clues about the nature of the engines powering these events. High-resolution, time-tagged data which makes a search for high frequency periodicities possible is available through the BATSE data archive at the Compton Observatory Science Support Center (COSSC).\\footnote[1]{% http://cossc.gsfc.nasa.gov/} We know of no other large-scale search for periodic emission that examines these data. ", + "conclusions": "\\label{sec:conc} This study is the first extensive search for high-frequency periodic emission to use the BATSE Time-Tagged Event data. The outcome of this study has not been the discovery of high-frequency periodicities in the burst records from SGRs and GRBs, but instead absence of evidence of these signals. In our study we have found just one plausible candidate for periodic emission at a frequency in our search range. The Rayleigh power obtained from trigger 7041 corresponds to detection near the 1\\% significance level, but the lack of emission at the same frequency in bursts from the same source (SGR~1900$+$14) at nearly the same time argues against this as a true detection. We have uncovered evidence for occasional TTE data corruption that produces a signal with $ f \\approx 900$~Hz; but the corrupt data is identifiable by obvious discrepancy between the TTE light curves and light curves using other data types. We have simulated burst data in an effort to determine how strong a periodic component must be in the signal from a burst to be detectable. A rough rule-of-thumb is that for a catalogue with $\\approx 2000$ bursts, periodic modulation comprising $\\approx 10$\\% of the photons in a given BATSE TTE record would have yielded significant Rayleigh power about half the time. (Since background is included in the record, the pulsed fraction in the burst signal must be higher.) Our analysis found oscillations at this level in only one TTE record, trigger 7041 from SGR 1900+14. That so few significant values of the Rayleigh power were found indicates that there is no substantial high-frequency periodic emission from these sources." + }, + "0112/astro-ph0112537_arXiv.txt": { + "abstract": "It has recently been argued that bubble nucleation in ekpyrotic and cyclic cosmological scenarios can lead to unacceptable inhomogeneities unless certain constraints are satisfied. In this paper we show that this is not the case. We find that bubble nucleation is completely negligible in realistic models. ", + "introduction": " ", + "conclusions": "" + }, + "0112/nucl-th0112003_arXiv.txt": { + "abstract": "In the present study we investigate the static properties of nuclei in the inner crust of neutron stars. Using the Hartree-Fock method in coordinate space, together with the semiclassical approximation, we examine the patterns of phase transitions. ", + "introduction": " ", + "conclusions": "" + }, + "0112/astro-ph0112054.txt": { + "abstract": "Whenever a freely spinning body is found in a complex rotational state, this means that either the body experienced some interaction within its relaxation-time span, or that it was recently ``prepared'' in a non-principal state. Both options are encountered in astronomy where a wobbling rotator is either a recent victim of an impact or a tidal interaction, or is a fragment of a disrupted progenitor. Another factor (relevant for comets) is outgassing. By now, the optical and radar observational programmes have disclosed that complex rotation is hardly a rare phenomenon among the small bodies. Due to impacts, tidal forces and outgassing, the asteroidal and cometary precession must be a generic phenomenon: while some rotators are in the state of visible tumbling, a much larger amount of objects must be performing narrow-cone precession not so easily observable from the Earth. The internal dissipation in a freely precessing top leads to relaxation (gradual damping of the precession) and sometimes to spontaneous changes in the rotation axis. Recently developed theory of dissipative precession of a rigid body reveals that this is a highly nonlinear process: while the body is precessing at an angular rate $ \\omega$, the precession-caused stresses and strains in the body contain components oscillating at other frequencies. Dependent upon the spin state, those frequencies may be higher or, most remarkably, lower than the precession rate. In many states dissipation at the harmonics is comparable to or even exceeds that at the principal frequency. For this and other reasons, in many spin states the damping of asteroidal and cometary wobble happens faster, by several orders, than believed previously. This makes it possible to measure the precession-damping rate. The narrowing of the precession cone through the period of about a year can be registered by the currently available spacecraft-based observational means. We propose an appropriate observational scheme that could be accomplished by comet and asteroid-aimed missions. Improved understanding of damping of excited rotation will directly enhance understanding of the current distribution of small-body spin states. It also will constrain the structure and composition of excited rotators. However, in the near-separatrix spin states a precessing rotator can considerably slow down its relaxation. This lingering effect is similar to the one discovered in 1968 by Russian spacecraft engineers who studied free wobble of a tank with viscous fuel. ", + "introduction": "In 1730, 23-year-old leutenant Leonhard Euler retired from the Russian navy, to become a professor of physics at the Russian Academy of Sciences. Eleven years later he interrupted his tenure, to assume the post of Director of Mathematics and Physics at the Berlin Academy, offered to him by King Friedrich the Second. There he stayed until 1765 when he was invited back to St.Petersburg by the enlightened Empress Catherine the Great. It is during that 25-year-long Berlin period of his life that Euler made his major contributions to the mechanics of a rotating body. After publishing some prefatory results in 1750 and 1758, Euler wrote down, in 1760, his now celebrated equations describing free spin of an unsupported top with arbitrary moments of inertia $\\;I_1,\\;I_2, \\;I_3\\;$: \\be I_i \\; {\\dot{\\Omega}}_i \\; - \\; \\left(I_j \\; - \\; I_k \\protect\\right) \\; \\Omega_j \\; \\Omega_k \\; = \\; \\tau_i \\; \\; \\; . \\label{1.1} \\ee $\\;\\Omega_{1,2,3}\\;$ being the the angular-velocity components in the coordinate system defined by the three principal axis of inertia, (1, 2, 3). For a freely rotating body, the external torques $\\;\\tau_{1,2,3}\\;$ standing in the right-hand side are nil. Without loss of generality, one may assume that $ \\;I_3\\,\\geq\\,I_2\\,\\geq\\,I_1\\;$, and thus, always to consider axis (3) to be the major-inertia axis. This result was published only five years later, in Chapter X of Euler's book {\\textit {Theoria motus corporum solidorum seu rigidorum ex primis nostrae cognitionis principiis stabilita et ad omnes motus qui in huiusmodi corpora cadere possunt accomodata.}} In Chapter XI, Euler easily found the solution for a prolate symmetrical case\\footnote{Later the symmetric top was also treated by Lagrange (1788) and Poisson (1813).}: if $\\;I_3\\,=\\,I_2\\,>\\,I_1\\;$ (like, for example, in the case of an elongated rod), the vector of inertial angular velocity $\\;\\bf \\Omega\\;$ describes a circular cone about the minor-inertia axis (1) of the body: \\be {\\Omega}_1 \\; \\; = \\; \\; const\\;\\;,\\;\\;\\; {\\Omega}_2 \\; \\; = \\; \\; {\\Omega}_{\\perp} \\cos {\\omega}t~~,~~~ {\\Omega}_3 \\; \\; = \\; \\; {\\Omega}_{\\perp} \\sin {\\omega}t~~,~~~ \\label{1.2} \\ee $\\omega =(I_1/I_3 - 1) \\Omega_1 $ being the precession rate. It is possible to show that the angular-momentum vector $\\,\\bf J\\,$ behaves like $\\,\\bf \\Omega\\,$, i.e., precesses about axis (1) at the same angular rate. This picture of $\\,\\bf \\Omega\\,$ and $\\,\\bf J\\,$ precessing about axis (1) is the one seen by an observer associated with the body frame. In an inertial observer's opinion, things will look different: from his viewpoint the angular momentum $\\,\\bf J\\,$ of an unsupported top will conserve, while its angular velocity $\\,\\bf \\Omega\\,$ and the minor-inertia axis (1) will be precessing about $\\,\\bf J$. Similarly, in the case of oblate symmetry ($\\;I_3\\,>\\,I_2\\,=\\,I_1\\;$) both $\\;\\bf \\Omega\\;$ and $\\;\\bf J\\;$ will perform, in the body frame, a circular precession about the major-inertia axis (3): \\be {\\Omega}_1 \\; \\; = \\; \\; {\\Omega}_{\\perp} \\cos {\\omega}t~~,~~~ {\\Omega}_2 \\; \\; = \\; \\; {\\Omega}_{\\perp} \\sin {\\omega}t~~,~~~ {\\Omega}_3 \\; \\; = \\; \\; const \\label{1.3} \\ee where $\\,\\omega = (I_3/I_1 - 1) \\Omega_3.\\;$ In an inertial frame, though, it will be $\\;\\bf \\Omega\\;$ and axis (3) describing circular cones about $\\;\\bf J\\;$ In Chapter XIII Euler tackled the general case of $I_3 >I_2 \\geq I_1$ and solved it in terms of functions presently known as elliptic integrals. These were pioneered a century earlier by John Wallis and Isaac Newton, and went in the late XVIII - early XIX centuries under the name of elliptic functions\\footnote{See, for example, http://www-groups.dcs.st-andrews.ac.uk/~history/HistTopics/}. (Euler, though, used neither of these names, and did not refer to Wallis or Newton.) Nowadays the name ``elliptic'' belongs to functions {\\textit {sn, cn, dn}} and their kin, that were not known at the time of Euler. They were introduced by Karl Jacobi in 1829 (Jacobi 1829), studied by Legendre (1837), and later employed (Jacobi 1849, 1882) in the rotating-top studies. These functions are, in a way, generalisations of our customary trigonometric functions: while for symmetric prolate and oblate bodies the circular precession is expressed by (\\ref{1.2}) and (\\ref{1.3}) correspondingly, in the general case $I_3 \\geq I_2 \\geq I_1$ the solution will read: \\begin{eqnarray} \\Omega_1\\;=\\;\\gamma\\;\\,{\\it{dn}}\\left(\\omega t , \\; k^2 \\protect\\right)\\;\\;, \\;\\;\\;\\;\\Omega_2 \\; = \\; \\beta \\, \\; sn\\left(\\omega t , \\; k^2 \\protect\\right) \\;\\;,\\;\\;\\;\\;\\Omega_3 \\; = \\; \\alpha\\;\\,{\\it{cn}}\\left(\\omega t,\\;k^2 \\protect\\right)\\;\\;,\\;\\;\\; \\label{1.4} \\end{eqnarray} for $\\;{\\bf{J}}^2 \\; < \\; 2\\;I_2 \\; T_{\\small{kin}} \\; $, and \\begin{eqnarray} \\Omega_1\\;=\\;{\\tilde \\gamma}\\;\\,{\\it{cn}}\\left({\\tilde \\omega} t,\\;{\\tilde k}^2 \\protect\\right)\\;\\;,\\;\\;\\;\\;\\Omega_2 \\; = \\;{\\tilde \\beta}\\,\\;{\\it sn}\\left( {\\tilde \\omega} t ,\\;{\\tilde k}^2 \\protect\\right)\\;\\;,\\;\\;\\;\\;\\Omega_3 \\; = \\; { \\alpha}\\;\\,{\\it{dn}}\\left({\\tilde \\omega} t,\\;{\\tilde k}^2 \\protect\\right) \\label{1.5} \\end{eqnarray} for $\\;\\;{\\bf{J}}^2\\;>\\;2\\;I_2\\;T_{\\small{kin}}\\;$. Similarity between (\\ref{1.2}) and (\\ref{1.4}), as well as between (\\ref{1.3}) and (\\ref{1.5}), is evident. In the above expressions, the precession rate $\\;\\it \\omega\\;$ and the parameters $\\;\\alpha ,\\;\\beta ,\\;{\\tilde \\beta}, \\;\\gamma ,\\;{\\tilde \\gamma}, \\;{\\tilde \\omega},\\;k\\;$ and $\\;{\\tilde k}\\;$ are some sophisticated combinations of $\\;I_{1,2,3}, \\;T_{\\small {kin}}\\;$ and $\\;{\\bf J}^2\\;$. What is important, is that solution (\\ref{1.4}) approaches (\\ref{1.2}) in the limit of prolate symmetry, $\\;(I_3\\,-\\,I_2)/I_1\\,\\rightarrow\\,0\\;$, while solution (\\ref{1.5}) approaches (\\ref{1.3}) in the limit of oblate symmetry, $\\;(I_2\\,-\\,I_1)/I_1\\,\\rightarrow\\,0\\;$. To adumbrate in an illustrative manner the applicability realms of (\\ref{1.4}) and (\\ref{1.5}), let us turn to Figure 1. In the course of free spin, two quantities (integrals of motion) are conserved. One is the angular momentum \\be {\\bf{J}}^2 \\;=\\;I_1^2\\,{\\Omega}_1^2\\;+\\;I_2^2\\,{\\Omega}_2^2\\;+\\;I_3^2\\, {\\Omega}_3^2 \\;\\;,\\;\\;\\; \\label{1.6} \\ee another is the kinetic energy \\be T_{\\small{kin}}\\;=\\;\\frac{1}{2}\\;\\left\\{I_1\\,{\\Omega}_1^2\\;+\\;I_2\\,{\\Omega}_2^2 \\;+\\;I_3\\,{\\Omega}_3^2 \\right\\}\\;\\;.\\;\\;\\; \\label{1.7} \\ee Evidently, these two expressions define ellipsoids in the angular-velocity space $\\;(\\Omega_1,\\,\\Omega_2,\\,\\Omega_3)\\;$. Intersection of these two surfaces will be the trajectory described by vector $\\;\\bf \\Omega\\;$ in the said space. On Figure 1, the angular-momentum ellipsoid is depicted. On its surface, we have marked the lines of its intersection with several different kinetic-energy ellipsoids appropriate to different energies. It does not take much space imagination to understand that, for a fixed angular-momentum surface (\\ref{1.6}), there exist an infinite family of kinetic-energy surfaces (\\ref{1.7}) intersecting with it. The largest surface of kinetic energy (corresponding to the maximal value of $\\;T_{kin}\\;$) is an ellipsoid that fully encloses our angular-momentum ellipsoid and only touches it in point A and its opposite. Similarly, the smallest surface of kinetic energy (corresponding to minimal $\\; T_{kin}\\;$) would be an ellipsoid fully contained inside our angular-momentum ellipsoid and only touching it from inside, at point C and its opposite. It is easy to demonstrate that, for a fixed $\\;\\bf J\\;$, the maximal and minimal possible values of the kinetic energy are achieved during rotations about the minimal-inertia and maximal-inertia axes, correspondingly. It can also be shown, from ($\\ref{1.1}$), that in the case of a non-dissipative and torque-free rotation, the tip of \\pagebreak \\break %%{\\bf Fig. 1} \\begin{figure}%%f1 \\centerline{\\epsfxsize=3.5in\\epsfbox{fig1.ps}} \\bigskip \\caption{The constant-angular-momentum ellipsoid, in the angular-velocity space. The lines on its surface are its intersections with the kinetic-energy ellipsoids corresponding to different values of the rotational energy. The quasi-stable pole A is the maximal-energy configuration, i.e., the state wherein the body spins about its minimal-inertia axis. The stable pole C symbolises the minimal-energy state, i.e., rotation about the maximal-inertia axis. The angular-velocity vector describes the constant-energy lines, and at the same time slowly shifts from one line to another, approaching pole C. The picture illustrates the case of an elongated body: $I_3 \\stackrel{>}{\\sim}I_2>I_1$. The trajectories are circular near A and remain (in the case of an elongated body) virtually circular almost up to the separatrix. The trajectories will regain a circular shape only in the closemost proximity of C.} \\end{figure} \\pagebreak \\noindent the vector $\\;{\\bf{\\Omega}}\\;$ will be describing, on Figure 1, a curve along which the angular-momentum and energy ellipsoids intersect (Lamy \\& Burns 1972). Hence, these intersections may be called trajectories. Solution ($\\ref{1.4}$) is valid for higher energies, i.e., from pole A through the separatrix; solution ($\\ref{1.5}$) works for lower energies, i.e., from the separatrix through pole C. Wherever the trajectories are almost circular, the solutions ($\\ref{1.4}$) and ($\\ref{1.5}$) may be approximated by ($\\ref{1.2}$) and ($\\ref{1.3}$), correspondingly. The formalism developed by Euler and refined by Jacobi might be a perfect tool for description of rotation of asteroids, comets, cosmic-dust granules, spacecrafts and whatever other unsupported rigid rotators, if not for one circumstance, inner dissipation. Because of this circumstance, the Euler-Jacobi theory of precession works only for time spans short enough to neglect dissipation. The presence of inner dissipation may be guessed even on a heuristic level. The bounded range of permissible energies makes one think that a freely spinning body of a fixed angular momentum must be seeking ways of relaxation, i.e., of getting rid of the excessive energy, in order to approach the minimal-energy configuration. Thence the necessity of some dissipation mechanism. Two such mechanisms are known. One is relevant only for mesoscopic rotators, like interstellar-dust grains, and therefore plays a certain role in the cosmic-dust alignment. This is the Barnett dissipation, a phenomenon called into being by the periodic remagnetisation of a precessing paramagnetic body (Lazarian \\& Draine 1997). The second mechanism, inelastic dissipation, is, too, relevant for mesoscopic grains (Lazarian \\& Efroimsky 1999), and it plays the decisive role in the macroscopic bodies' relaxation. The effect results from the alternating stresses produced inside a wobbling body by the time-dependent acceleration of its parts. The stresses deform the body, and the inelastic effects cause dissipation of the rotational energy. The dissipation entails relaxation of the precession: the major-inertia axis of the body and its angular-velocity vector $\\;\\bf \\Omega \\;$ tend to align along the angular momentum $\\bf J$. In other words, the precession cone described by $\\;\\bf \\Omega \\;$ about $\\bf J$ will be narrowing until $\\;\\bf \\Omega \\;$ aligns along $\\bf J$ completely. A simple calculation (Efroimsky 2001, Efroimsky 2000, Efroimsky \\& Lazarian 2000, Lazarian \\& Efroimsky 1999) shows that in this case the major-inertia axis of the body will align in the same direction, so that, from the body-frame viewpoint, $\\;\\bf \\Omega \\;$ will eventually be pointing along this axis. This configuration will correspond to the minimal kinetic energy, the angular momentum being fixed. An inertial observer will thus see the unsupported body miraculously changing its rotation axis. This is exactly what happened in 1958 when, to mission experts' surprise, satellite Explorer I changed its rotation axis. The spacecraft was a very elongated body. It had been supposed to spin about its least-inertia axis (i.e., about its longest dimension), but refused to do so, and instead started precessing (Thomson 1961). This was probably the first example of a practical need for a further development of the Eulerian theory of a free top, a development that would address an unsupported top with dissipation. Another motivation for this work was put forward in the same year by Prendergast (1958) who studied the asteroid population of the Solar System and enquired as to how many asteroids could be in non-principal (i.e., precessing) spin states, and how this could evidence of the impact frequency in the main belt. (Prendergast implied that it is collisions that might drive asteroids out of the principal state and make them wobble.) An important point made by Prendergast was the generation of the second harmonic: if a body is precessing at an angular rate $\\;\\omega\\;$, then the dissipation is taking place not only at this frequency but also at double thereof. Prendergast failed to notice the emergence of the higher harmonics, but even his noticing of the second harmonic was an important observation. In several other aspects the mathematical treatment of the problem, offered by Prendergast, was erroneous and gave him no chance to come to a reasonable solution. Moreover, at that time the observational astronomy lacked any reliable data on wobbling asteroids. So, Prendergast's paper was forgotten (even though once in a while it appeared in the references), and his successors had to start up from scratch. The interest in the asteroidal precession re-emerged in 70-s, after the publication of the important work (Burns \\& Safronov 1973) that suggested estimates for the relaxation time, based on the decomposition of the deformation pattern into bulge flexing and bending, and also on the conjecture that ``the centrifugal bulge and its associated strains wobble back and forth relative to the body as the rotation axis {\\bf $\\; \\bf \\omega\\;$} moves through the body during a wobble period.'' As turned out later, the latter conjecture does not work, because the inelastic dissipation, for the most part of it, is taking place not near the surface but in the depth of the body, i.e., not right under the bulge but deep beneath it. Thus, the bulge is much like an iceberg tip. This became clear when the distribution of precession-caused stresses was calculated, with improved boundary conditions (Efroimsky \\& Lazarian 2000), (Lazarian \\& Efroimsky 1999). Another, main, problem of Burns \\& Safronov's treatment was their neglection of the nonlinearity, i.e., of the second and higher harmonics. The nonlinearity, in fact, is essential. Neglection thereof leads to a large underestimation of the damping rate, because the leading effect comes often from the second and higher harmonics (Efroimsky \\& Lazarian 2000), (Efroimsky 2000). All in all, the neglection of nonlinearity and mishandling of the boundary conditions leads to a several-order underestimate of the precession-damping rate. In the same year, Peale published an article dealing with inelastic relaxation of nearly spherical bodies (Peale 1973), and there he did take the second harmonic into account. In 1979 Purcell addressed a similar problem of interstellar-grain precession damping. He ignored the harmonics and mishandled the boundary conditions upon stresses: in (Purcell 1979) the normal stresses had their maximal values on the free surfaces and vanished in the centre of the body (instead of being maximal in the centre and vanishing on the surfaces). These oversights lead to a several-order underevaluation of the dissipation effectiveness and, thereby, of the relaxation rate. ", + "conclusions": "We may be very close to observation of the relaxational dynamics of wobbling small Solar System bodies, dynamics that may say a lot about their structure and composition and also about their recent histories of impacts and tidal interactions. Monitoring of a wobbling comet during about a year after it leaves the 3 AU zone will, most probably, enable us to register its precession relaxation. %\\textsf 1. Dissipation at the second and higher harmonics makes a major input into inelastic-relaxation process. Distribution of stresses and strains over the volume of a precessing body is such that an overwhelming share of the inelastic dissipation is taking place deep inside the body, not in its shallow regions, as thought previously. These and other reasons make the inelastic relaxation several order more effective than believed hitherto. 2. A finite resolution of radar-generated images puts a limit on our ability of recognising as to whether an object is precessing or not. Relaxation-caused changes of the precession-cone half-angle may be experimentally observed. Our estimates show that the modern spacecraft-based instruments are well fit for observations of the asteroid and cometary wobble relaxation. Relaxation may be registered within relatively short periods of time (about a year). 3. Measurements of the relaxation rate will provide us with a valuable information on attenuation in asteroid and cometary materials, as well as on their recent history of impacts and tidal interactions. 4. Since the inelastic relaxation is far more effective than presumed earlier, the number of asteroids expected to wobble with a recognisable half-angle of the precession cone must be lower than expected. (We mean the predictions suggested in (Harris 1994).)" + }, + "0112/astro-ph0112270_arXiv.txt": { + "abstract": "Current evolutionary models imply that most Cataclysmic Variables (CVs) have $P_{\\rm orb}<2 $ hours and are Dwarf Nova (DN) systems that are quiescent most of the time. Observations of nearby quiescent DN find that the UV spectrum is dominated by the hot white dwarf (WD), indicating that it provides a significant fraction of the optical light in addition to the quiescent disk and main sequence companion. Hence, identifying a faint, quiescent CV in either the field or a globular cluster (GC) from broadband colors depends on our ability to predict the WD contribution in quiescence. We are undertaking a theoretical study of the compressional heating of WDs, extending down to very low time averaged accretion rates, $\\timav\\sim 10^{-11}M_\\odot \\ {\\rm yr^{-1}}$, which allows us to self-consistently find the $T_{\\rm eff}$ of the WD. We demonstrate here that most of the compressional heating occurs in the freshly accreted envelope and that the WD core temperature reaches a fixed value on a timescale less than typical evolutionary times. Since nuclear burning is unstable at these $\\timav$'s, we have incorporated the recurrent heating and cooling of the WD core throughout the classical novae limit cycle in order to find the $T_{\\rm eff}$-$\\timav$ relations. Comparing to observations of field DN confirms the $\\timav$-$P_{\\rm orb}$ relation of disrupted magnetic braking. We also predict broad-band colors of a quiescent CV as a function of $\\timav$ and companion mass and show that this leads to the identification of what may be many CVs in deep HST images of GCs. ", + "introduction": "Dwarf Novae (DN) systems contain a white dwarf (WD) accreting matter at time-averaged rates $\\timav<10^{-9}M_\\odot \\ {\\rm yr}^{-1}$ from a low-mass ($<0.5M_\\odot$ typically) stellar companion (see Osaki 1996 for an overview). At these $\\timav$'s, the accretion disk is subject to a thermal instability which causes it to rapidly transfer matter onto the WD (at $\\dot M \\gg \\timav$) for a few days to a week once every month to year. The orbital periods of these binaries are usually less than 2 hours (below the period gap), but there are also DN above the period gap, $>$ 3 hours (see Shafter 1992). The $\\dot M$ onto the WD is often low enough between outbursts that the UV emission is dominated by the internal luminosity of the WD. Indeed recent spectroscopy has resolved the WD's contribution to the quiescent light and found effective temperatures $T_{\\rm eff}\\sim 10,000-40,000 \\ {\\rm K}$ (see Sion 1999). The measured internal WD luminosity is larger than expected from an isolated WD of similar age ($\\approx$ Gyr), indicating that it has been heated by accretion (Sion 1985). Compressional heating (i.e. internal gravitational energy release) appears to be the main driver for this re-heating (Sion 1995). Sion's (1995) estimate for internal gravitational energy release within the WD (of mass $M$ and radius $R$) was $L\\approx 0.15 GM\\timav/R$. However, we show in \\S \\ref{sec:physics} that the energy release actually depends on the thermal state of the WD interior and that the dominant energy release is in the accreted outer envelope, giving $L\\approx 3kT_c \\timav/\\mu m_p$, where $\\mu\\approx 0.6$ is the mean molecular weight of the accreted material, $T_c$ is the WD core temperature, $m_p$ is the baryon mass, and $k$ is Boltzmann's constant. The theoretical challenge that we address in \\S \\ref{sec:Tcmethod} is how to calculate $T_c$ as a function of $\\timav$, and thus find $T_{\\rm eff}$. Because of unstable nuclear burning and the resulting classical novae cycle, the H/He envelope mass changes with time, allowing the core to cool at low accumulated masses and be heated prior to unstable ignition. We use nova ignition to determine the maximum mass of the overlying freshly accreted shell, and find the steady-state (i.e. cooling equals heating throughout the classical novae cycle) core temperature, $T_c$, as a function of $\\timav$ and $M$. We compare our calculations to HST/STIS observations and infer $\\timav$ on the timescale of $10^6$ years. We find that DN above the period gap have $\\timav\\approx 10^{-9}M_\\odot \\ {\\rm yr^{-1}}$, while those below have $\\timav\\approx 10^{-10} M_\\odot \\ {\\rm yr^{-1}}$, consistent with that expected from traditional CV evolution (e.g. Howell, Nelson, \\& Rappaport 2001), even those that involve some ``hibernation'' (Shara et al.\\ 1986; Kolb et al.\\ 2001). The result is more surprising if the much weaker magnetic braking laws of Andronov, Pinsonneault, \\& Sills (2001) are correct. We also predict the minimum light ($M_V$) of CVs in quiescence for a range of $\\timav$, WD mass, and companion mass. This assists the search for the predicted large population of CVs with very low mass companions ($<0.1M_\\odot$) that are near, or past, the period minimum (Howell, Rappaport, \\& Politano 1997). Observations already show that the WD fixes the quiescent colors of these CVs and our calculations are useful for CV surveys in the field (e.g. 2DF, SDSS, see Marsh et al.\\ 2001 and Szkody et al.\\ 2002) and globular clusters. ", + "conclusions": "We have evaluated the action of compressional heating on accreting WD interiors and shown that most of the compressional energy release takes place in the accreted envelope, and is thermally communicated to the core. The maximum envelope mass is set by the unstable nuclear burning that causes a classical nova runaway and most likely expels the accreted mass. We have constructed equilibrium accretors which have constant core temperatures such that the heat lost from the core when the envelope is thin (i.e. right after the classical nova) is balanced by that regained when the envelope is thick. This equilibrium determines the $T_{\\rm eff}$ of the WD throughout the classical nova cycle. Our models agree with the observations of dwarf novae in deep quiescence and imply $\\timav\\approx10^{-10}M_\\odot$ yr$^{-1}$ just below the period gap and $\\timav\\approx10^{-9}M_\\odot$ yr$^{-1}$ just above the period gap for WD masses in the range $0.6$--$1.0M_\\odot$. Our $T_{\\rm eff}$ calculations provide a prediction of the colors of quiescent DN. Using MS stellar models, we have predicted where a DN should appear in a color-magnitude diagram as a function of $\\timav$ and the mass of its companion. Many unidentified objects appear in the relevant regions of the detailed CMDs which have been obtained for globular clusters by HST. The number of such systems in the field will increase due to upcoming surveys (such as SDSS and 2DF, see Marsh et al.\\ 2001), and will push to lower $\\timav$ systems. Though our initial efforts have met with apparent success, there is still much to be done. We need to vary the metallicity of the accreted material, lowering to values appropriate for globular cluster science. This could change our results at large $\\timav$, but at low $\\timav$'s, the ignition mass is set by $pp$ burning and will likely not change too much. We also need to relax our initial assumptions, e.g. by including WD excavation or accretion and accounting for thermal evolution of the WD. The internal thermal state of the WD has been a longstanding uncertainty in classical nova work, as has the question of how much mass is ejected in the explosion (Gehrz et al.\\ 1998). Our work provides the first calculation of the internal thermal state of a WD undergoing classical novae, and will eventually lead to self consistent calculations for ignition masses, including variations of the metallicity. This will be an improvement on previous work (e.g. Prialnik \\& Kovetz 1995) which treated $T_c$ and $\\timav$ as two independent parameters." + }, + "0112/astro-ph0112046_arXiv.txt": { + "abstract": "The theory of shock acceleration predicts the maximum particle energy to be limited only by the acceleration time and the size (geometry) of the shock. This led to optimistic estimates for the galactic cosmic ray energy achievable in the SNR shocks. The estimates imply that the accelerated particles, while making \\emph{no strong impact on the shock structure} (test particle approach) are nevertheless scattered by \\emph{strong self-generated} Alfven waves (turbulent boost) needed to accelerate them quickly. We demonstrate that these two assumptions are in conflict when applied to SNRs of the age required for cosmic ray acceleration to the {}``knee{}'' energy. We study the \\emph{combined} effect of acceleration nonlinearity (shock modification by accelerated particles) and wave generation on the acceleration process. We show that the refraction of self-generated waves resulting from the deceleration of the plasma flow by the pressure of energetic particles causes enhanced losses of these particles. This effect slows down the acceleration and changes the shape of particle spectrum near the cut-off. The implications for observations of TeV emission from SNR remnants are also discussed. ", + "introduction": "The first-order Fermi or diffusive shock acceleration (DSA) has been long considered as responsible for the production of galactic cosmic rays (CRs) in supernova remnants (SNRs), as well as for the radio, \\( x \\)- and \\( \\gamma \\)-ray emission from these and other shock related objects. The most crucial characteristic of this process that is usually examined in terms of its capability to explain a given observation, is the rate at which it operates. Indeed, what is often expected from the theory or even inferred from the observations is an extended particle energy spectrum, frequently a power-law, but more rapidly decaying at the highest energies observed. Often, this decay is referred to as an energy or momentum cut-off and is usually associated with the finite acceleration time or with losses if their rate exceeds the acceleration rate. As long as the losses are unimportant, the cut-off \\( p_{max}(t) \\) advances with time according to the following equation \\begin{equation} \\label{p:max} \\frac{dp_{max}}{dt}=\\frac{p_{max}}{t_{acc}} \\end{equation} whereas in the presence of losses the acceleration rate \\( p_{max}/t_{acc} \\) may be equated to the loss rate to yield a steady state value of \\( p_{max} \\). The acceleration time scale is determined by (\\eg \\citealt{axf81}) \\begin{equation} \\label{t:acc} t_{acc}=\\frac{3}{u_{1}-u_{2}}\\int _{p_{min}}^{p_{max}}\\left[ \\frac{\\kappa _{1}(p)}{u_{1}}+ \\frac{\\kappa _{2}(p)}{u_{2}}\\right] \\frac{dp}{p} \\end{equation} with \\( u_{1} \\) and \\( u_{2} \\) being the upstream and downstream flow speeds in the shock frame and with \\( \\kappa _{1,2} \\) being the particle diffusivities in the respective media. One may recognize in the last formula the sum of average residence times of a particle spent upstream and downstream of the shock front before it completes one acceleration cycle, integrated over the entire acceleration history from \\( p_{min} \\) to \\( p_{max} \\). Given the flow speeds \\( u_{1,2} \\) which, in many cases are known reasonably well, the most sensitive quantity is the particle diffusivity \\( \\kappa \\). This, in turn, is determined by the rate at which particles are pitch angle scattered by the Alfven turbulence. If the latter was just a background turbulence in the interstellar medium (ISM), the acceleration process would be too slow to produce the galactic CRs in SNRs (\\eg \\citealt{lc83}). However it was realized (\\eg \\citealt{bell78a,bla:ost}) that accelerated particles should create the scattering environment by themselves generating Alfven waves on the cyclotron resonance \\( kp\\mu /m=\\omega _{ci} \\), where \\( k \\) is the magnitude of the wave vector (directed along the magnetic field), \\( p \\), \\( \\mu \\), \\( m \\) and \\( \\omega _{ci} \\) are the particle momentum, the cosine of its pitch angle, mass and non-relativistic (\\( eB/mc \\)) gyro-frequency. Note, that the diffusive character of particle transport (and determination of \\( \\kappa \\)) has been rigorously obtained within a quasi-linear theory, \\ie it is subject to constraints on the turbulence level. The wave generation, however, proved to be very efficient (see \\eg \\citealt{vdm84} and the next section). In particular, using, again, the quasi-linear approximation, the normalized wave energy density \\( \\left( \\delta B/B_{0}\\right) ^{2} \\) may be related to the partial pressure \\( P_{c} \\) of CRs that resonantly drive these waves through \\begin{equation} \\label{delB} \\left( \\delta B/B_{0}\\right) ^{2}\\sim M_{A}P_{c}/\\rho u^{2} \\end{equation} where \\( M_{A} \\) is the Alfven Mach number and \\( \\rho u^{2} \\) is the shock ram pressure. Since \\( M_{A} \\) is typically a large parameter, \\( \\delta B/B_{0} \\) may become larger than unity even if the acceleration itself is relatively inefficient, \\ie if \\( P_{c}/\\rho u^{2}\\ll 1 \\). Strictly speaking this invalidates the quasi-linear approach as a means for describing the generation of strong turbulence at shocks. The commonly accepted way to circumvent this difficulty is to assume that the turbulence saturates at \\( \\delta B/B_{0}\\sim 1 \\), which means that the m.f.p. of pitch angle scattered particles is of the order of their gyro-radius \\( r_{g} \\). Then, \\( \\kappa =\\kappa _{B}\\equiv c \\)\\( r_{g}(p)/3 \\), where the speed of light \\( c \\) is substituted instead of CR velocity and \\( \\kappa _{B} \\) stands for the Bohm diffusion coefficient. This immediately sets the acceleration time scale (\\ref{t:acc}) at the level of particle gyro-period \\( \\left( eB/p\\right) ^{-1} \\) times \\( \\left( c/u_{1}\\right) ^{2} \\). In principle, the turbulence level \\( \\delta B/B_{0} \\) significantly exceeding unity is possible in local shock environments (see \\eg numerical studies by \\citealt{be95} and \\citealt{bell:luc}). As a consequence of that the diffusion coefficient could be even smaller than \\( \\kappa _{B} \\), and hence, the acceleration rate would be faster than it is commonly believed to be. At the same time, since usually Alfvenic type turbulence is considered, the respective velocity perturbations must be super-Alfvenic and supersonic, which raises questions about its ability to sustain itself in an extended area without rapid dissipation that will decrease the \\( \\delta B/B_{0} \\) level. Likewise, decreasing of turbulence level below the Bohm limit, for example due to the finite extent of the turbulence zone upstream, should slow down the acceleration \\citep{lc83}. However the acceleration rate given by eq.(\\ref{t:acc}) with \\( \\kappa =\\kappa _{B} \\) was found to be fast enough to explain (at least marginally) the acceleration of CRs in SNRs up to the {}``knee{}'' energy \\( \\sim 10^{15}eV \\) over their life time. Much further optimism has been caused by the studies of \\citet{dav} and \\citet{naito}. They analyzed the prospects for detection of super-TeV emission from nearby SNR that should be produced by the decays of \\( \\pi ^{0} \\) mesons born in collisions of shock accelerated protons with the nuclei of interstellar gas. The expected fluxes were shown to be detectable by the imaging \\( \\rm \\check{C} \\)erenkov telescopes. Moreover, the EGRET \\citep{esp} detected a lower energy (\\( \\la GeV \\)) emission coinciding with some galactic SNRs. The spectra also seemed consistent with the DSA predictions. One may even argue that the low energy EGRET data verified one of the most difficult elements of the entire acceleration mechanism, the so called injection. In essence, this is a selection process (not completely understood) whereby a small number of thermal particles become subject to further acceleration (see \\citealt{gjk00,zank01} for the latest development of the injection theory and \\citealt{mdru01} for a review) and may be then treated by standard means of the DSA theory that was designed to describe particles with velocities much higher than the shock velocity. Therefore, what seemed left for the theory was to continue the EGRET spectrum (that sets the normalization constant, or injection rate) with some standard DSA slope (nearly \\( E^{-2} \\) or somewhat steeper) and to predict the \\( \\gamma \\)-ray flux in the TeV range where it could be detected by the \\( \\rm \\check{C} \\)erenkov telescopes. Unfortunately, despite the physical robustness of the arguments given by \\cite{dav,naito}, no statistically significant signal that could be attributed to any of the EGRET sources was detected. The further complication is that some critical energy band between GeV and TeV energies is currently uncovered by available instruments. Therefore, based on these observational results it was suggested (\\eg \\citealt{buck98}) that there is probably a spectral break or even cutoff somewhere within this band. However the spectrum above GeV energies remains an enigma. This will be resolved perhaps with the launch of the GLAST mission and when the new generation of \\( \\rm \\check{C} \\)erenkov telescopes with lower energy thresholds begin to operate. However, the discovery of the 100 TeV emission from SNR1006 \\citep{tanim98}, as well as some other remnants not seen by the EGRET at lower energies (see, \\eg \\citealt{ahar01,allen01,kirkd01} for a complete discussion), although almost universally identified with electrons diffusively accelerated to similar energies, is widely interpreted as a strong support of the mechanism itself. The above suggests, however, that in reality it might be not as robust as is its simplified test particle version with enhanced turbulence and particle scattering. In this paper we attempt to understand what may happen to the spectrum provided that the acceleration is indeed fast enough to access the TeV energies over the life time of SNRs in question. Our starting point is that the fast acceleration also means that the pressure of accelerated particles becomes significant in an early stage of Supernova evolution so that the shock structure is highly nonlinear. At the first glance this should not slow down acceleration since, according to eq.(\\ref{delB}), this changes the turbulence \\emph{level} thus improving particle confinement near the shock front and thus making acceleration faster (smaller \\( \\kappa \\)). However, the formation of a long CR precursor (in which the upstream flow is gradually decelerated by the pressure of CRs, \\( P_{c} \\)) influences the \\emph{spectral properties} of the turbulence by affecting the propagation and excitation of the Alfven waves. This effect is twofold. First the waves are compressed in the converging plasma flow upstream and are thus blue-shifted, eliminating the long waves needed to keep exactly the highest energy particles diffusively bound to the accelerator. Second, and as a result of the first, at highest energies there remain fewer particles than expected so that the level of resonant waves is smaller and hence the acceleration rate is lower. We believe that these effects have been largely overlooked before which may have substantially overestimated the particle maximum energy in strongly nonlinear regimes. ", + "conclusions": "There are at least two reasons to believe that the standard acceleration theory may have estimated the maximum particle energy or the form of the spectrum below it incorrectly. The first reason is simply a possible conflict with the observations of TeV-emission from SNRs as we discussed in the Introduction section. The second reason is a theoretical one, that arises naturally from considering the nonlinear response of the shock structure to the acceleration which is exemplified in Figure~\\ref{fig:bif}. According to this picture, the response is so strong that it is unlikely that the acceleration can proceed at the same rate with no change in physics after such a dramatic shock restructuring (pre-compression \\( R \\) may rise by 1-2 orders of magnitudes depending on the Mach number). Time dependent numerical simulations (\\eg \\citealt{kjg02}) show that the modifications occur very quickly, and compression is increased substantially even before \\( p_{max}\\sim 1 \\) (note that this would be consistent with the bifurcation diagram in Figure~\\ref{fig:bif} for initial \\( \\nu \\sim (c/u_{1})n_{CR}/n_{1}\\ga 0.1 \\), where \\( n_{CR}/n_{1} \\) is the ratio of CR number density at the shock to that of the background plasma far upstream). The shock modification, in turn, must follow rather abruptly after the maximum momentum has passed through the critical value. It was argued recently (\\citealt{mdv00}) that this should drive crucial acceleration parameters such as the maximum momentum and injection rate back to their critical values which must limit shock modification and settle it at some marginal level, the so called self-organized critical (SOC) level (see also \\citealt{tom00, mdru01, mdFer01} for more discussions of the critical interrelation between the injection, maximum energy and shock structure). Mathematically, the SOC state is characterized by the requirement of merging of the two critical points on the bifurcation diagram in Figure~\\ref{fig:bif} into one inflection point on the \\( \\nu (R) \\) graph. Perhaps the most appealing aspect of this approach is its ability to predict the values of all three order parameters (injection rate, maximum momentum and compression ratio) given the only control parameter (the Mach number) just from our knowledge of the nonlinear response \\( R(\\nu ,p_{max}) \\) shown in Figure~\\ref{fig:bif}, and no further calculations. However, the required backreaction mechanisms on the injection and maximum momentum need to be demonstrated to operate. We have already discussed at a qualitative level how the injection rate is reduced by shock modification. The subject of this paper has been the reduction of particle momenta related to the formation of a spectral break at \\( p=p_{*} \\), as a result of wave compression in a modified shock precursor. The position of the spectral break is universally related to the degree of system nonlinearity \\( R \\), since \\( p_{*}=p_{max}/R \\). Hence, the problem seems to be converted to the study of nonlinear properties of the acceleration that are formally known from the analytic solution shown in Figure~\\ref{fig:bif}. However, the injection rate \\( \\nu \\) that is now required for accurate determination of the spectral break \\( p_{*} \\) through \\( R \\), may currently be obtained only from the SOC ansatz. It should be also mentioned that strong reduction of \\( p_{*} \\) is obviously not to be expected in oblique shocks, where the resonance relation \\( kp\\propto B \\) is approximately preserved due to the compression of \\( B \\) simultaneously with \\( k \\). An equally important problem is that strong losses of particles between \\( p_{*} \\) and \\( p_{max} \\) must slow down the growth of \\( p_{max}(t) \\) due to the reduction of resonant waves. As we argued in section~\\ref{sec:signif}, this may result in an order of magnitude slower acceleration than one would expect from the standard Bohm diffusion paradigm. Consequently the dynamically and observationally significant spectral break \\( p_{*} \\) may be at least two orders of magnitude below the maximum momentum \\( p_{max} \\) (again, depending on \\( M \\)) that could be reached in the unimpeded acceleration which is normally implied in estimates of maximum energy achievable in SNRs over their active life time. In addition to the above mentioned uncertainty in \\( p_{max}(t) \\), its relation to the position of the spectral break \\( p_{*}=p_{max}/R \\) also needs further clarification. Indeed, since \\( R \\) depends on a dynamical cut-off \\( \\hat{p} \\) which, in general, is linked to \\( p_{*} \\) and \\( p_{max} \\), the latter relation is still implicit. It can be easily resolved, however, in a supercritical regime (the saturated part of the \\( R(\\nu ) \\) dependence in Figure~\\ref{fig:bif}, see also \\citealt{mdru01} for details), which requires\\footnote{% This is strictly valid for \\( \\kappa (p)\\propto p \\). } \\( \\nu \\hat{p}/p_{inj}\\gg M^{3/4} \\). One simply has then \\( R\\approx M^{3/4} \\). As it was argued, however, the injection is unlikely to be high enough to reach this regime. An additional argument against it is that the spectral break becomes unrealistically small in the \\( M\\rightarrow \\infty \\) limit, since \\( p_{*}=p_{max}/M^{3/4} \\). In the opposite case \\( \\nu \\hat{p}/p_{inj}\\ll M^{3/4} \\), the compression rate saturates at \\( R\\approx \\nu \\hat{p}/p_{inj} \\). Note that the injection rate must be still above critical, otherwise \\( R\\approx 1 \\). Now we need to specify \\( \\hat{p} \\). The simple approximation used in the previous section yielded \\( \\hat{p}\\approx 0.1p_{max} \\), so that \\( p_{*}\\approx 10p_{inj}/\\nu \\) (independent of \\( p_{max} \\)) which may be regarded as a lower bound on \\( p_{*} \\). Indeed, the above relation between \\( \\hat{p} \\) and \\( p_{*} \\) may be applied only to the outermost part of the shock transition (see Secs. \\ref{sec: ext},\\ref{sec:connect}). Downstream, the spectrum cuts off very sharply immediately beyond \\( p_{*} \\), section~\\ref{sec: int}. Therefore, the dynamical cut-off \\( \\hat{p}\\approx p_{*} \\) and we obtain the following upper bound on \\( p_{*} \\), \\( p_{*}\\approx \\sqrt{p_{inj}p_{max}(t)/\\nu } \\). It should be clear that unless \\( \\nu \\) is dramatically reduced as a result of shock modification, even this upper bound places \\( p_{*} \\) way below \\( p_{max} \\). This may be the reason for non-detection of protons at TeV energies in SNRs. Finally, this does not contradict to the detection of 10-100 TeV electrons in \\eg SNR 1006 since they may be accelerated by other mechanisms (\\eg \\citealt{pap81, gal84, bykuv99, lam01}) or may have higher radiation efficiency." + }, + "0112/astro-ph0112100_arXiv.txt": { + "abstract": "We present the source catalogue for the SCUBA Lens Survey. We summarise the results of extensive multi-wavelength observations of the 15 submillimetre-selected galaxies in the catalogue, from X-rays to radio. We discuss the main observational characteristics of faint submillimetre galaxies as a population, and consider their interpretation within the framework of our understanding of galaxy formation and evolution. ", + "introduction": "The highly successful far-infrared(IR) all-sky survey undertaken by {\\it IRAS} led to the identification of numerous highly obscured star-forming and active galaxies in the local Universe, $z\\ls 0.3$ (Soifer, Neugebauer \\& Houck 1987). These systems are some of the most luminous galaxies at the present day and emit most of their radiation in the far-IR waveband, although they contribute only 0.3 per cent of the local luminosity density (Sanders \\& Mirabel 1996). More recent work in the far-IR and submillimetre (submm) wavebands has produced a similar revolution in our view of obscured galaxies in the {\\it distant} $z\\gs 1$ Universe. These observations have employed the {\\it COBE} and {\\it ISO} satellites and the Submm Common-User Bolometer Array (SCUBA; Holland et al.\\ 1999) on the 15-m James Clerk Maxwell Telescope\\footnote{The JCMT is operated by the Joint Astronomy Centre on behalf of the United Kingdom Particle Physics and Astronomy Research Council (PPARC), the Netherlands Organisation for Scientific Research, and the National Research Council of Canada.} (JCMT). The new observations have shown that the ultraluminous far-IR population evolves more strongly than the equivalent optically-selected population and that, in contrast to the local Universe, luminous obscured galaxies at high redshift could contribute a substantial fraction of the total emitted radiation. This conclusion is confirmed by comparing the energy density in the optical (Bernstein et al.\\ 2001) and far-IR/submm backgrounds (Puget et al.\\ 1996; Fixsen et al.\\ 1998; Finkbeiner, Schlegel \\& Davis 2000). These backgrounds represent the cumulative energy emitted in these wavebands across all epochs, mainly at redshift $z \\sim 1$. The approximate equivalence of the energy density in the two regimes shows that somewhere near half of the total radiation in the Universe came from obscured energy sources, which could be either stars or AGN. If the majority of this emission is powered by radiation from stars with a standard initial mass function (IMF), then approximately half of all the stars that have formed by the present day could have formed in highly obscured systems. Clearly it is critically important to include these highly-obscured sources in models of galaxy evolution if we are to obtain a complete understanding of the formation and evolution of galaxies. The advent of sensitive submm imaging with SCUBA has allowed a number of groups to undertake surveys for distant submm galaxies. Results on the number density of sources in blank fields as a function of 850-$\\mu$m flux density have been published by three groups: Hughes et al.\\ (1998) worked with a single deep map centred on the {\\it Hubble Deep Field} (HDF), while Barger et al.\\ (1998, 1999b) employed a combination of deep/narrow and wide/shallow observations of fields in the Lockman Hole and Hawaii Survey Field regions; finally, there has been a survey of areas from the Canada-France Redshift Survey (Eales et al.\\ 1999, 2000; Lilly et al.\\ 1999). Shallower, wider surveys have also been carried out by Borys et al.\\ (2001) and by the UK Submm Survey consortium (Dunlop 2001; Scott et al.\\ 2001; Fox et al.\\ 2001; Almaini et al.\\ 2001). Due to the modest resolution of SCUBA, 15$''$ FWHM at 850\\,$\\mu$m, the deepest of these studies are confusion limited at $\\sim 2$\\,mJy (Hughes et al.\\ 1998); this is the deepest flux density for which reliable source detection is possible in blank fields (Blain, Ivison \\& Smail 1998; Hogg 2001). At slightly longer wavelengths, observations using the MAMBO 1.2-mm camera on the IRAM 30-m telescope have recently been reported in both field and lensing cluster regions (Bertoldi et al.\\ 2000). Rather than mapping large fields, another approach to construct large samples of submm-detected galaxies is to exploit the tight correlation between the far-infrared and radio luminosities of star-forming galaxies (Condon 1992) and use SCUBA to target samples of faint radio sources (Chapman et al.\\ 2001b, 2001c). This innovative technique is particularly well-suited for identifying the wide-field samples needed to tackle issues such as the clustering strength of SCUBA galaxies (e.g.\\ Almaini et al.\\ 2001). Our collaboration has adopted yet another approach with the aim of pushing below the confusion limit of the blank-field surveys. We achieve this by using massive gravitational cluster lenses to increase both the sensitivity and resolution of SCUBA (Blain 1998). The first submm counts were based on maps of two clusters (Smail, Ivison \\& Blain 1997). The survey was subsequently expanded to cover seven lensing clusters at $z=0.19$--0.41 (Smail et al.\\ 1998). The results of similar surveys have recently been reported by Chapman et al.\\ (2001) and van der Werf et al.\\ (2001b). These observations of lensing clusters benefit from a typical amplification factor of 2--3$\\times$, improving both the sensitivity of the maps and their effective resolution, and so allowing confusion-free counts to be derived down to $\\sim 0.5$\\,mJy (Blain et al.\\ 1999a), well below the conventional 2-mJy field confusion limit for SCUBA. In fortuitous cases, the amplification can exceed $10\\times$ (e.g.\\ van der Werf et al.\\ 2001b), providing the opportunity to identify submm galaxies as faint as $\\sim 0.1$\\,mJy and study their properties. Based upon this survey we have published the number counts of submm galaxies (Smail, Ivison \\& Blain 1997; Blain et al.\\ 1999a), the identification of the counterparts to the submm sources in the optical (Smail et al.\\ 1998), near-infrared (Smail et al.\\ 1999a; Frayer et al.\\ 2000), radio (Smail et al.\\ 2000) and X-ray (Fabian et al.\\ 2000) bands, as well as optical spectroscopy of candidate counterparts (Barger et al.\\ 1999a). We have also provided detailed multi-wavelength follow-up observations of the brighter sources (Ivison et al.\\ 1998a, 2000a). Building on the redshifts determined for the three brightest sources in the sample, we have obtained the first CO detections of submm-selected galaxies using the OVRO and IRAM interferometers (Frayer et al.\\ 1998, 1999; Kneib et al.\\ 2001). Yet higher-resolution CO images of one source have been obtained by combining data from the OVRO and BIMA arrays (Ivison et al.\\ 2001). High-resolution mm-continuum observations using the OVRO array have also been presented (Frayer et al.\\ 2000). Similar mm-continuum observations of sources in other SCUBA surveys have been presented by Downes et al.\\ (1999), Gear et al.\\ (2000) and Lutz et al.\\ (2001). Finally, the interpretation of these observations and their relevance to our understanding of galaxy formation and evolution at high redshifts has been discussed by Blain et al.\\ (1999b, 1999c). This paper includes a summary and update of these previous results. As a benchmark for the following discussion we note that an Ultraluminous Infrared Galaxy (ULIRG) with a far-IR luminosity of $L_{\\rm FIR} \\sim 3\\times 10^{12} L_\\odot$ (similar to Arp\\,220), and thus a star-formation rate (SFR) of $\\sim 300$\\,M$_\\odot$\\,yr$^{-1}$ would have a 850-$\\mu$m flux density of $\\gs 3$\\,mJy out to $z\\sim 10$ in a Universe with $q_0=0.5$.\\footnote{We assume $q_o=0.5$ and $H_o=50$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$ throughout the paper, except in Fig.~1.} The effects of different SEDs and world models on the results in the two SCUBA observing bands are illustrated in Fig.~1. In three 8-hr shifts of observing in good conditions with SCUBA it is possible to achieve a 3$\\sigma$ flux limit of 3\\,mJy at 850\\,$\\mu$m across a 160$''$-diameter field (upgrades to SCUBA now mean that this limit is reached in closer to two shifts), probing a volume of 10$^6$ Mpc$^3$ out to $z\\sim 10$ for ULIRGs. \\begin{figure} \\centerline{\\psfig{file=f1.ps,angle=270,width=3.3in} } \\noindent{\\small\\addtolength{\\baselineskip}{-3pt}} \\caption{The flux density at 450 and 850\\,$\\mu$m expected from a ULIRG-like galaxy as a function of redshift for different dust temperatures and cosmologies. The galaxy is assumed to have $L_{\\rm FIR} = 3 \\times 10^{12}$\\,L$_\\odot$.} \\end{figure} Here we present the source catalogue of luminous submm galaxies from our SCUBA survey, and discuss the current identifications and properties of their counterparts in other wavebands. We then summarise the broad characteristics of the populations that contribute to the submm background radiation. In \\S2 we briefly describe the submm observations and their reduction and analysis. In \\S3 we present our updated estimate of the source counts at 850 and 450\\,$\\mu$m, and in \\S4 we describe the identification of counterparts and constraints on their likely redshifts and power sources. A brief summary of the characteristics of each source is given in \\S5, and we discuss the properties of the different classes of submm galaxy in \\S6. In \\S7 we present the main conclusions arising from the SCUBA lens survey and summarise these in \\S8. ", + "conclusions": "In summary, the bulk of the background radiation intensity has been resolved into discrete submm sources using SCUBA at an 850-$\\mu$m intrinsic flux limit of $\\sim 1$\\,mJy. The counterparts to these submm sources appear to be dusty, ultraluminous galaxies with very diverse optical and near-IR properties. The majority of these galaxies appear to lie at $z\\gs 1$, with a median redshift of $z\\sim 2.5$--3. The optically brighter sources have been studied in more detail using mm-wave interferometers, and they show the large dynamical masses and high gas fractions expected for young massive galaxies. The characteristics of these galaxies are consistent with them being the progenitors of the most massive elliptical galaxies seen in the local Universe. A comparison of the detailed properties of a handful of these galaxies with local ULIRGs suggests that the rapid increase in dust-obscured activity at high redshifts, need to explain the submm counts and the FIRB, has its origin in the increasing instability of the gas-rich, bulge-weak progenitors of the submm population at high redshifts. However, the majority of the mJy submm population remain elusive; they have very faint (or invisible) counterparts in the optical and near-IR and progress in investigating their nature and properties is likely to be slow. In the future we look forward to increases in the number of submm galaxies with accurate redshift determinations -- crucial for follow-up CO line-mapping to provide dynamical masses and gas fractions. Work in this area will require deep spectroscopy in the near-IR (and optical) on 10-m class telescopes for the brighter Class-I and II sources, as well as more innovative approaches, such as blind radio searches for OH/H$_2$O maser emission (Townsend et al.\\ 2001). Improved constraints on the redshift distribution of the whole submm population await the confirmation of the radio-submm spectral index (and more detailed SED fitting) as a reliable estimator of redshift for submm-selected galaxies. In part this will rely on checking at higher luminosities and redshifts, using {\\it SIRTF} and SOFIA, the relatively weak dependence of dust temperature on luminosity seen in low-redshift {\\it IRAS}-selected galaxies by Dunne et al.\\ (2000). Equally essential is the detailed study of the characteristics of individual sources. Here the main advances are likely to come from observations across a wide range of wavelength, with the X-ray, mid-IR and far-IR wavebands being the most promising, providing crucial information about the distribution of dust temperatures and the power sources driving these systems. Observations at higher spatial resolution, in the near-IR, mm and radio wavebands, will allow us to study the internal structure of these galaxies -- to search for morphological evidence of the events which triggered their prodigous activity. In particular, the refurbishment of {\\it NICMOS} on-board {\\it HST} will provide a powerful tool to interpret the rest-frame optical morphologies of these galaxies and compare them to local ULIRGs to test if the same physical processes are responsible for ultraluminous systems at low and high redshifts. In the longer term, the direct study of the detailed astrophysics of submm galaxies will benefit immensely from the sensitivity and resolving power of the 10-milliarcsec resolution ALMA interferometer array. We expect that submm surveys which exploit lens amplification, such as the one presented here, will retain a central role in studying the submm galaxy population. In part this is because the properties of these galaxies tax the capabilities of current instrumentation in many wavebands and hence the boost provided by the lens is essential for successful follow-up. Moreover, observations through massive gravitational lenses allow us to probe intrinsically fainter submm sources, which are more representative of the population responsible for the bulk of the FIRB. We look forward to continued exploitation of the sample presented here, and the results of new surveys, to study the nature of the faint submm population and answer some of the questions raised in this paper. Finally, the goal of future theoretical work in this area should be to incorporate both the obscured submm population (Class~0 and I) and the less-obscured systems (Class-II submm galaxies and the more massive classical LBGs) into a single evolutionary sequence and hence naturally explain the relation between the two populations. Important observational input on this question can be obtained by studying the relative clustering of the various populations in well-defined environments at high redshifts, in particular the overdense regions around some luminous radio galaxies (Ivison et al.\\ 2000b) may evolve into the cores of massive clusters at the present day. Such studies will require wide-field surveys covering the UV, near-IR and longer wavelengths, the new SCUBA2 submm camera for the JCMT and the WFCAM panoramic near-IR camera for UKIRT will be a powerful facilities for obtaining the essential observations in these wavebands." + }, + "0112/astro-ph0112385_arXiv.txt": { + "abstract": "{Recent X-ray observations reveal an increasing number of X-ray sources in nearby galaxies exceeding luminosities of $L_{\\rm x}\\ga2 \\cdot 10^{39} \\mbox{erg}\\ \\mbox{s}^{-1}$. Assuming isotropic emission, the Eddington limit suggests a population of intermediate-mass black holes of $M_\\bullet\\gg10M_\\odot$. However, \\citeN{MarkoffFalckeFender2001} proposed that jets may be contributing to the X-ray emission from X-ray binaries (XRBs), implying that some X-ray sources may be relativistically beamed. This could reduce the required black hole masses to standard values. To test this hypothesis, we investigate a simple X-ray population synthesis model for X-ray point sources in galaxies with relativistic beaming and compare it with an isotropic emission model. The model is used to explain a combined data set of X-ray point sources in nearby galaxies. We show that the current distributions are consistent with black hole masses $M_\\odot\\la10$ and bulk Lorentz factors for jets in microquasars of $\\gamma_{\\rm j}\\sim5$. Alternatively, intermediate mass black holes up to 1000 $M_\\odot$ are required which are distributed in a powerlaw with roughly $\\frac{dN}{dM} \\sim M^{-2}$. ", + "introduction": "In recent years X-ray observations have revealed several off-nucleus ultra-luminous X-ray sources (ULXs) in the luminosity range $10^{39}\\ - \\ 10^{40}\\mbox{erg}\\ \\mbox{s}^{-1}$ within nearby galaxies (\\citeNP{LaParolaPeresFabbiano2001}; \\citeNP{MizunoKubotaMakishim2001}; \\citeNP{BauerBrandtSambruna2001}; \\citeNP{ColbertMushotzk1999}). The Eddington limit for an accreting object with mass $M$ is $L_{\\rm Edd} \\approx 1.25 \\cdot 10^{38} \\frac{M}{M_\\odot} \\mbox{erg}\\ \\mbox{s}^{-1}$, which implies that these sources are super-Eddington for stellar mass objects. Some ULXs show spectral transitions from a soft spectrum to a hard power law and luminosity variability (e.g. \\citeNP{MizunoKubotaMakishim2001}; \\citeNP{KubotaMizunoMakishima2001}), ruling out supernova remnants and supporting the idea that ULXs can be attributed to accreting black holes. To achieve the observed X-ray luminosities with isotropically radiating accretion disks should require a population of intermediate-mass black holes of $50-500 M_\\odot$. As discussed in \\citeN{KubotaMizunoMakishima2001}, however, the measured inner-disk temperatures of ULXs ($T_{\\rm{in}} = 1.0-1.8$ keV) are too high for these masses and there is no established formation scenario for such intermediate-mass black holes. The problems with isotropic emission models have already been discussed by \\citeN{KingDaviesWard2001}, where the authors propose some form of anisotropic emission as an alternative. A beaming factor of ten already reduces the required mass of the black holes to expected values, but this is difficult to achieve with pure disk models. Recently \\citeN{MarkoffFalckeFender2001} suggested that the spectrum of some XRBs could be explained by a coupled disk/jet model, where some of the X-ray emission is produced by synchrotron and inverse-Compton radiation in a jet. This emission would naturally be relativistically beamed. \\citeN{MirabelRodriguez1999} (see also \\citeNP{ReynoldsLoanFabian1997}) have pointed out that a number of nearby galaxies should host microblazars - microquasars with relativistically beamed jets pointed towards the observer. We will here investigate whether such populations of microblazars or intermediate mass black holes can indeed explain current data on ULXs and constrain the basic parameters required for these models. ", + "conclusions": "Using two very simple models for the evolution of XRBs, we calculate the luminosity distribution of X-ray point sources in nearby galaxies. We consider a jet/disk model based on \\citeN{FalckeBiermann1999} and \\citeN{MarkoffFalckeFender2001}, which can give rise to relativistically beamed emission from microblazars. Alternatively we also consider a purely isotropically radiating disk model. Both models can in principle reproduce a combined luminosity function compiled from X-ray point source catalogs of three close galaxies and the XHFS spiral galaxy sample. However, as expected, the isotropic disk model requires a mass distribution of black holes extending out to $1000 M_\\odot$ to explain the ULXs. On the other hand, a relativistic jet/disk model can fit the data with stellar mass black holes, if X-ray emitting jets with Lorentz factors $\\gamma_{\\rm j}\\simeq5$ are present in XRBs. In addition, a fraction of $\\eta=10-30\\%$ of the total soft X-ray emission has to come from the jet rather than the accretion disk for an un-beamed XRB in the high state. This requires rather powerful jets but is not completely unreasonable. If only a fraction of the XRBs have relativistic jets, a slightly higher Lorentz factor or jet efficiency is needed. Boosting a $10$ mJy Galactic XRB by a factor $\\sim 10^2$ (for $\\gamma \\sim 5$) and placing it at $D \\sim 3 {\\rm Mpc}$ would yield only a faint $10$ nJy source and make radio detections difficult. With the current statistics it is not possible to distinguish between the two different models, but it seems that microblazars provide at least a sensible alternative to the often discussed intermediate mass black hole scenario. Monitoring the spectral variability of the most luminous sources and further developing the XRB jet model should eventually help to disentangle the two scenarios. \\vspace{3mm} \\noindent {\\it Acknowledgments} The authors would like to thank R. Fender and an anonymous referee for useful suggestions." + }, + "0112/astro-ph0112450_arXiv.txt": { + "abstract": "{ In this paper we present some results concerning the effects of two instantaneous starbursts, separated by a quiescent period, on the dynamical and chemical evolution of blue compact dwarf galaxies. In particular, we compare the model results to the galaxy IZw18, which is a very metal-poor, gas-rich dwarf galaxy, possibly experiencing its first or second burst of star formation. We follow the evolution of a first weak burst of star formation followed by a second more intense one occurring after several hundreds million years. We find that a galactic wind develops only during the second burst and that metals produced in the burst are preferentially lost relative to the hydrogen gas. We predict the evolution of several chemical abundances (H, He, C, N, O, $\\alpha$-elements, Fe) in the gas inside and outside the galaxy, by taking into account in detail the chemical and energetical contributions from type II and Ia supernovae. We find that the abundances predicted for the star forming region are in good agreement with the HII region abundances derived for IZw18. We also predict the abundances of C, N and O expected for the HI gas to be compared with future FUSE abundance determinations. We conclude that IZw18 must have experienced two bursts of star formation, one occurred $\\sim 300$ Myr ago and a present one with an age between 4-7 Myr. However, by taking into account also other independent estimates, such as the color-magnitude diagram and the spectral energy distribution of stars in IZw18, and the fact that real starbursts are not instantaneous, we suggest that it is more likely that the burst age is between 4 and 15 Myr. ", + "introduction": "Blue Compact Dwarf galaxies (BCD) are characterized by compact appearance, high gas content, very blue colors and low chemical abundances. These properties are typical of unevolved systems, thus suggesting that BCD should have suffered very few bursts of star formation during their lives and that some of them are probably experiencing their first burst (Searle \\& Sargent 1972). In a recent review, Kunth \\& \\\"Ostlin (2000) argued that, despite a few remaining young galaxy candidates (like IZw18, SBS0335-052 or HS0822+3542; Lipovetsky et al. 1999; Kniazev et al. 2000), in most BCDs an old underlying stellar population does exist, revealing at least another burst of star formation (SF) besides the present one. In a recent survey of BCD, \\\"Ostlin et al. (2001) have found that, although the young burst population dominates the integrated optical luminosities, it contributes only marginally to the total stellar mass. It is possible that each starburst episode in BCD be followed by a quiescent (or almost quiescent) period, with a time scale of the order 10$^8$ to 10$^9$ yr (Leitherer 2001), during which winds produced by supernova explosions expel the gas out of the region and the star formation fades. When the luminosity of the burst is no more able to sustain the wind, the gas should cool and collapse back toward the center of the galaxy, thus allowing for the onset a new burst (Babul \\& Rees 1992; D'Ercole \\& Brighenti 1999). Off-centered SN explosions can also drive inward-propagating shocks, thus creating the conditions of a new SF event in the center of the galaxy (Mori et al. 2001). The galaxy IZw18 is the most metal-poor local galaxy known so far and was considered until recently as the best candidate for a truly ``young'' galaxy. Stellar population analyses by Hunter \\& Thronson (1995) and Dufour et al. (1996) were not deep enough to reveal any old stellar population, but recent studies of deeper Color-Magnitude Diagrams (CMD), both in the optical (Aloisi, Tosi \\& Greggio 1999; hereafter ATG) and in the infrared (\\\"Ostlin 2000) revealed the presence of two stellar populations in IZw18: a young population with an age of $\\sim 15$ Myr and an asymptotic giant branch population with an age of several $10^{8}$ years. Chemical evolution models by Kunth et al. (1995) fit the abundances observed in IZw18 with one, or at maximum two short bursts. Evolutionary population synthesis models by Mas-Hesse \\& Kunth (1999; hereafter MHK) showed that the present burst is very young (ranging between 3 and 13 Myr, depending on whether the burst is instantaneous or continuous) and the contribution of the stars of an ancient burst to the emission over the whole UV-optical range, if any, is negligible. Legrand (2000) and Legrand et al. (2000) proposed instead a low and continuous SF regime for IZw18; furthermore they assumed that the observed metals cannot result from the material ejected by the aging starburst, because these metals are hidden in a hot phase and therefore undetectable when using the optical spectroscopy. In a previous paper (Recchi, Matteucci \\& D'Ercole 2001; hereafter Paper I) we presented a study concerning the effect of a single, instantaneous starburst on the dynamical and chemical evolution of a gas-rich dwarf galaxy, with galactic parameters resembling those of IZw18. We showed that the observed abundances of IZw18 could be reproduced also by a single-burst model provided that its age is $\\sim$ 30 Myr, occurring in a primordial gas (zero initial metallicity), although we did not exclude the presence of an underlying older population which polluted only slightly the ISM (less than 1/100 of solar metallicity). However, the estimate of $\\sim 30$ Myr for IZw18 seems to be in contrast both with the spectral energy distribution and the color-magnitude diagram, as mentioned above. Other results of Paper I can be summarized as follows: \\par i) a galactic wind develops as a consequence of the starburst and carries away mostly the metals produced during the starburst. In particular, we found that the metals produced by type Ia SNe are lost even more efficiently than the metals produced by type II SNe. This fact is important since different SN types produce different elements, in particular, the net effect is to enhance the $\\alpha$-element over Fe ratio inside the galaxy relative to the gas lost from the galaxy. \\par ii) The cooling of metals in the gas was found to be very efficient so that most of the metals should be found in the cold gas phase already after few Myr from the beginning of the burst. Both results i) and ii) depend on the assumed energy transferred from SNe into the interstellar medium (ISM), a crucial parameter in galaxy evolution studies. \\par In this paper we intend to test the hypothesis of multiple starbursts with our model. In order to investigate this hypothesis, we performed some numerical simulations to study in detail the dynamical and chemical evolution of a gas-rich dwarf galaxy experiencing two single, instantaneous starbursts. Our approach makes use of a two-dimensional dynamical model coupled with detailed chemical enrichment from both type II and type Ia SNe. In section 2 we summarize the properties of IZw18, in section 3 we describe the model and the assumptions adopted in our simulation. The results are presented and discussed in section 4. Finally in section 5 some conclusions are drawn. ", + "conclusions": "In this paper we have studied the chemo-dynamical evolution of the gas as a consequence of two instantaneous starburts, separated by a quiescent period, in a galaxy similar to IZw18. We have taken into account in detail both SN (Ia and II) feed-back and stellar yields. The main conclusions can be summarized as follows: \\begin{itemize} \\item in most of the explored cases the wind develops only during the second burst and the newly formed metals are ejected more efficiently than the pristine gas (mostly H), thus confirming previous results (Paper I; D'Ercole \\& Brighenti 1999; MacLow \\& Ferrara 1999). \\item In particular, the metals produced by type Ia SNe are lost more efficiently than those produced by type II SNe, due to the higher efficiency in energy transfer by type Ia SNe, which explode in an already hot and rarefied medium. However this effect, already found in Paper I where only one burst was assumed, is milder in the case of 2 bursts, since SNeII in the second burst are also likely to transfer more energy into the ISM. \\item Our best model suggests that a first weak burst of star formation, occurred roughly 300 Myr ago, followed by a more intense one, can well reproduce the properties of IZw18. \\item The predicted chemical abundances for the gas which remains inside the star-forming region are in very good agreement with the HII region abundances of IZw18, especially for what concerns the C/O and N/O ratios. The agreement with the data of IZw18 is definitively better for the case with 2 bursts than with only one, thus confirming the presence of old underlying stellar populations (ATG, \\\"Ostlin 2000). These results are in good agreement with observations either when using the chemical yields from low and intermediate mass stars by Renzini \\& Voli (1981) or those by van den Hoek \\& Groenewegen (1997). \\item In the framework of the standard nucleosynthesis in massive stars we find two solutions for the age of the second burst: 1) an age of some tens of Myr (depending on the nucleosynthetic prescriptions adopted) and 2) an age of $4-6$ Myr. This second solution is in good agreement with the spectral energy distribution models studied by MHK. If we assume that massive stars produce only primary N, the situation changes since the right N/O ratio is achieved already after 3-4 Myr after the first burst. Therefore, in this case we cannot use the N/O as a cosmic clock. However, the assumption of massive stars producing only primary N is perhaps too extreme. Therefore, we conclude that on the basis of only chemo-dynamical results we cannot safely suggest the age of the present time burst. As a consequence of this, we checked our burst age by computing the spectral energy distribution expected for the second burst, in the best model case. To do this we have computed $U-B$ and $B-V$ colors by adopting the package Starburst99 (Leitherer et al. 1999), a web based software and data package designed to model spectrophotometric and related properties of star-forming galaxies. The results obtained with Starburst99 were then compared with the observed values of IZw18 (namely $U-B = -0.88 \\pm 0.06$ and $B-V = -0.03 \\pm 0.04$; van Zee et al. 1998) and they are shown in Fig. 11. From this figure one can see that the outputs of Starburst99 are consistent with the observed $U-B$ only if the second burst has an age between 5.3 and 11 Myr, whereas the $B-V$ color is well reproduced if the age is between 23 and 44 Myr. However, the $B-V$ color is affected by the presence of the older stellar population (impossible to simulate with the Starburst99 package), whereas the $U-B$ color is dominated by young stars. Thus, these results seem to indicate that a very young age for the second burst (of the order of 5 -- 6 Myr) is the most likely solution, at least in the hypothesis of an instantaneous burst of star formation. Moreover, we can estimate the age of the second burst also from dynamical considerations. In particular, in our best model the bubble has travelled a distance of 720 pc in only 6-7 Myr from the beginning of the starburst. This distance is exactly the space between the shell and the center of the NW HII region observed in IZw18 (Martin, 1996). Martin (1996) herself suggests an age larger than the one derived here (see Table 4). This difference is due to the fact that Martin (1996) assumed an ISM distribution unmodified by the previous burst. In fact, the bubble expansion timescale strongly depends on the gas distribution. In conclusion, our results combined with the information we have from other studies, as summarized in Table 4, and with the fact that real starbursts are not instantaneous, suggest that one can reasonably conclude that the age of the present burst should be in the range 5-15 Myr. \\item Models with a flat ($x=0.5$) IMF produce much more oxygen after the first burst of star formation compared to models with the Salpeter (1955) IMF. For these models the only possible solution is a second burst of a very short age (around 4 Myr). \\item The [$\\alpha$/Fe] ratios in the two burst case are always lower in the gas lost through the wind than in the gas which remains bound to the galaxy, thus creating an asymmetry. However this effect, already present in the one burst case, is here less strong. The [$\\alpha$/Fe] inside the galaxy is predicted to be higher than solar (overabundance of $\\alpha$-elements) during the first 30 Myr from the burst. In particular, we predict a value of [O/Fe]$\\sim$+0.4 dex when the standard yields are adopted and a value of [O/Fe] $\\sim +0.7-0.8$ dex when the Fe in massive stars is artificially lowered by a factor of two, since the Woosley \\& Weaver (1995) yields for Fe seem to be overestimated. The last value of [O/Fe] is in very good agreement with the abundance analysis of IZw18 by Levshakov et al. (2001). \\item Finally, we computed the abundances of metals which should pertain to the HI gas and estimated that they are roughly a factor of 2 lower than the abundances in the star forming region (HII), a prediction which should be tested on new FUSE data. \\end{itemize} \\begin{table} \\begin{centering} \\caption[]{Ages of the second burst from various sources} \\begin{tabular}{ccc} \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} age (Myr) & source & reference \\\\ \\noalign{\\smallskip} \\hline\\noalign{\\smallskip} 5-6 & model M300 & this paper \\\\ 40-70 & model M300 & this paper \\\\ $\\sim$ 4 & model M300F & this paper \\\\ $\\sim$ 6 & model M500 & this paper \\\\ 40-120 & model M500 & this paper \\\\ 6-7 & bubble dynamics (M300) & this paper\\\\ \\hline 5-11 & U-B & S99 \\\\ 23-44 & B-V & S99 \\\\ 13.6 & $M_B$ & S99 \\\\ \\hline 15-20 & optical CMD & ATG \\\\ 15-20 & NIR CMD & Ostlin 2000\\\\ 15-27 & bubble dynamics & Martin 1996 \\\\ 3-13 & integrated spectra & MHK \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\vspace{0.3cm} Notes: We quote as S99 our own application of the Starburst99 (Leitherer et al. 1999) models to our examined cases. See text for details. \\end{centering} \\end{table} \\begin{figure} \\centering \\vspace{1.4cm} \\epsfig{file=recchi_fig11.eps, height=7.5cm,width=9cm} \\caption[]{\\label{fig:fig 11} {$U-B$ and $B-V$ colors computed by means of Starburst99 (Leitherer et al. 1999), for the model M300 (solid line) and M300F (dashed line), compared with published observed values for IZw18, with relative error-bars (dashed areas).}} \\end{figure} Before concluding we want to address some comments about the limits of the present model. In particular, the above results refer to the case of two subsequent instantaneous bursts of SF, the most common scenario usually attributed to the evolution of BCD. Until recently BCDs were supposed to undergo short and intense bursts of star formation, separated by long quiescent intervals, whereas dwarf irregulars and giant irregulars appeared to have a more continuous SF activity (e.g. Tosi 1999). For this reason, the vast majority of chemical evolution models for BCDs (e.g. Matteucci \\& Tosi 1985, Pilyugin 1993, Bradamante, Matteucci, D'Ercole 1998) were computed assuming burst durations of 10$^8$ years or less. Nowadays, however, there is increasing evidence that also BCDs have rather a {\\it gasping} SF, with long episodes of activity, separated by short quiescent intervals, if any. This continuity has been hypothesized by a few authors (e.g. Carigi et al. 1995, Legrand 2000), and has been supported by observations when the Hubble Space Telescope has allowed to resolve faint single stars even in galaxies outside the Local Group. The application of the method of synthetic color-magnitude diagrams (CMD) to derive the SF history in relatively distant BCD has provided fundamental information on the evolution of these systems. Deep CMD have been obtained from HST optical and infrared photometry of several BCD by different groups (e.g. Aloisi et al. 1999, Schulte-Ladbeck et al. 2001, Tosi et al. 2001) and have provided similar scenarios for the SF histories of the examined objects: the SF activity has started long ago (in all cases, at least as long ago as given by the maximum lookback time corresponding to the depth of the available photometry), but has been very intense only in a few galaxies. Standard episodes have a duration of several 10$^8$ years and tend to overlap each other with no real quiescent phases in between, or fairly short ones (lasting only 5--10 Myr). If the latter scenario turns out to apply in general to most BCD, the next fundamental step for a better understanding of their chemical evolution will be to take into account the effect on the ISM of SN explosions produced during long-lasting SF episodes. In fact, in most of the model solutions selected here, the predicted N abundances cross the observed range in a very short time interval (c.f. Figs. 5 and 7). We expect that models with continuous SF will predict abundances remaining within the observed range for longer, more realistic timescales. Despite the much heavier computer requirements, we are currently working on this implementation in our treatment and the results will be presented in a forthcoming paper." + }, + "0112/astro-ph0112499_arXiv.txt": { + "abstract": "The huge worldwide investment in CMB experiments should make the Sunyaev-Zeldovich (SZ) effect a key probe of the cosmic web in the near future. For the promise to be realized, substantial development of simulation and analysis tools to relate observation to theory is needed. The high nonlinearity and dissipative/feedback gas physics lead to highly non-Gaussian patterns that are much more difficult to analyze than Gaussian primary anisotropies for which the procedures are reasonably well developed. Historical forecasts for what CMB experiments might see used semi-analytic tools, including large scale map constructions, with localized and simplified pressure structures distributed on a point process of (clustered) sources. Hydro studies beyond individual cluster/supercluster systems were inadequate, but now large-volume simulations with high resolution are beginning to shift the balance. We illustrate this by applying ``Gasoline'' (parallelized Tree+SPH) computations to construct SZ maps and derive statistical measures. We believe rapid Monte Carlo simulations using parameterized templates centered on point processes informed by optical and other means on the observational side, and by hydro simulations on the theory side, should play an important role in pipelines to analyze the new SZ field data. We show that localized sources should dominate upcoming SZ experiments, identify sources in the maps under filtering and noise levels expected for these experiments, use the RCS photometric optical survey as an example of redshift localization, and discuss whether cosmic web patterns such as superclusters can be enhanced when such extra source information is supplied. ", + "introduction": "Hydro Approaches to SZ Forecasts} \\noindent {\\bf 1.1 The Resolution of Upcoming SZ Experiments:} The Compton upscattering of CMB photons by hot inhomogeneous (nonlinear) gas leads to secondary CMB anisotropies $\\Delta T/T (\\hq ,\\nu) =-2y_C (\\hq )\\psi_K(h\\nu /k_BT_\\gamma )$ in direction $\\hq$ at frequency $\\nu$, where $y_C = (\\sigma_T /m_ec^2) \\int n_e k_B (T_e-T_\\gamma)\\, d{\\tt l.o.s} $ is the Compton y-parameter and $\\psi_K$ depends only upon frequency.\\footnote{Here, $\\sigma_T$ is Thompson cross section, $m_e$ is the electron mass, $T_\\gamma$ is the CMB photon temperature, and $d{\\tt l.o.s}$ is the $\\hq$-line-of-sight radial distance element. $\\psi_K(x)=2-(x/2)(e^x+1)/(e^x-1)$ is 1 in the Rayleigh-Jeans region, zero at 217 GHz, negative above.} It is because $y_C$ is a direct probe of the line integral of the electron pressure in the hot intergalactic medium that the heavy investment in SZ experiments is so worthwhile. SZ observations of individual rich clusters have been possible for a decade, are now routine (\\eg Carlstrom \\etal 1999), and complementary to X-ray, optical and weak lensing observations. For resolved sources, the surface brightness of an SZ source is independent of its redshift. Even with the expected source evolution, this property should make the SZ effect a valuable probe of the cluster/group near/mid-field even at redshifts $z \\sim 1$, when we expect the system to be in a very active merging state. Thus, the era of blank field (or ambient) SZ surveys is upon us, some targeting the $\\sim 1^\\prime$ resolution well-matched to the cluster/group system at $z \\sim 1$, others with $\\sim 5^\\prime$ resolution, probing larger sky fractions, albeit with considerably enhanced source confusion. Specifications of a sample of upcoming/proposed SZ experiments are: \\ms \\noindent {\\bf Bolometer-based:} ACBAR: Viper telescope, 16 element, multi-frequency, $\\sim 4^\\prime$ {\\it fwhm} resolution, now; Bolocam/CSO: 10.4m CSO telescope, 151 pixels, 150, 220, 270 GHz, $1^\\prime$, fall 2001; Bolocam on the LMT?; ACT: 3 32x32 pixel bolometer array, $1.7^\\prime$, proposed; Planck: $\\sim 7^\\prime$ at 150 GHz, $\\sim 5^\\prime$ at 220 GHz, full sky, 2007. \\ms \\noindent {\\bf HEMT-based Interferometers:} OVRO mm array: 6 dish, 10.4m, 30 GHz, now; BIMA: 10 dish, 6.1m, 30 GHz, now; CARMA: OVRO+BIMA; CBI: 13 dish, 0.9m, 30 GHz, $\\sim 4^\\prime$; SZA: 6 dish, 3.5m, 30 GHz (+90 GHz), $\\sim 2^\\prime$; AMIBA: 19 dish, 1.2m+0.3m, 90 GHz, $\\sim 2^\\prime$; AMI: 10+8 dish, 3.7m+13m, 15 GHz, $\\sim 2^\\prime$. The resolution can be improved by spreading the dishes to longer baselines. \\vs \\noindent {\\bf 1.2 Historical Semi-analytic SZ Forecasting:} SZ estimates and limits derived from them were influential ever since Sunyaev and Zeldovich proposed the effect. First, attention was paid to baryon-dominated (BDM) models, with entropy injection via shocks or radiation, then to shock-heated neutrino-dominated (HDM) models, and then the many variants of cold dark matter (CDM) models, often with strong ``feedback'' of energy into the pregalactic or intergalactic medium. For example, in the eighties one of us (Bond \\etal 1980s) did HDM forecasts, first using the popular pancake treatment of structure formation, then a better cluster-based treatment using density peaks of various masses, including both Poisson and continuous clustering contributions of these ``shots'' (B88). Explosion-dominated models of structure formation in CDM, BDM models, from very massive objects, galaxies, superconducting strings, extreme preheating, \\etc \\, were also addressed in B88. All of these models became severely challenged by the COBE/FIRAS data which gave a $10^{-4}$ 95\\% CL upper bound to the allowed fraction of the CMB radiation in a Compton cooling distortion. Early SZ maps using cluster/group-scale peaks in CDM models were also constructed in the eighties (B90). In the nineties, the calculations of SZ maps became more sophisticated with the peak-patch technique, which included correct spatial clustering of halos, but ``painted on'' simple parameterized pressure-profiles within the halos (Bond \\etal 1990s, Bond \\& Myers 1996, Bond \\& Crittenden 2001 [BC01]). \\vs \\noindent {\\bf 1.3 Hydro Approaches:} Early single-cluster SPH calculations (Bond \\etal 1990s) became more sophisticated as the codes and computing power improved, as described for example in the ITP ``adiabatic'' cluster comparison test (Frenk \\etal 2000), for which SZ, X and weak lensing maps are shown in BC01. Ideas of the cosmic web interconnections of clusters of peak-patches were used to create optimally-designed rare supercluster (TreePM-SPH) simulations (Bond \\etal 1998, reviewed in BC01) that included a proper tidal field acting on its 104 Mpc high resolution patch (comoving lattice spacing $1.0\\mpc$, best $z$=0 resolution $20\\kpc$) and cooling (but no feedback) to see which, if any, of the X, SZ or weak lensing probes could be sensitive to mid-field and far-field structures around clusters and groups (\\eg the far-field filamentary bridges). As expected, we found weak lensing probes the far-field better than SZ, which does better than X. However, a critical issue is how to reveal such extended-source patterns, given the projected contributions from the clusters and environs behind and ahead of the supercluster targets (Sec.~3). In recent years, simulated SZ maps generated from hydrodynamical cosmological simulations have been used to predict what the experiments should see, mostly with emphasis on low order statistics such as the angular power spectrum. Even with the great improvements in computing power and codes we have seen, only a relatively small number of hydro realizations per SZ map-making exercise are being done, and so all such work remains statistically incorrect. Nonetheless, we believe these approximate treatments are useful steps along the path to that brave day when the full redshift range relevant to the projected maps from, say, 0 to 2, is do-able, the space tiled by contiguous simulation patches self-consistently constructed to have coherent long-waves joining them, all at the required resolution. This is routinely and rapidly done with the painted-halo peak-patch approach (Bond \\& Myers 1996). ", + "conclusions": "" + }, + "0112/astro-ph0112516_arXiv.txt": { + "abstract": "We present near-infrared spectroscopic observations of SN~1987A covering the period 1358 to 3158 days post-explosion. This is the first time that IR spectra of a supernova have been obtained to such late epochs. The spectra comprise emission from both the ejecta and the bright, ring-shaped circumstellar medium (CSM). The most prominent CSM emission lines are recombination lines of H~I and He~I, and forbidden lines of [S~III] and [Fe~II]. The ejecta spectra include allowed lines of H~I, He~I and Na~I and forbidden lines of [Si~I], [Fe~I], [Fe~II], and possibly [S~I]. The intensity ratios and widths of the H~I ejecta lines are consistent with a low-temperature Case~B recombination spectrum arising from non-thermal ionisation/excitation in an extended, adiabatically cooled H-envelope, as predicted by several authors. The slow decline of the ejecta forbidden lines, especially those of [Si~I], indicates that pure non-thermal excitation was taking place, driven increasingly by the decay of $^{44}$Ti. The ejecta iron exhibits particularly high velocities (4000--4500~km/s), supporting scenarios where fast radioactive nickel is created and ejected just after the core-bounce. In addition, the ejecta lines continue to exhibit blueshifts with values $\\sim$--200~km/s to --800~km/s to at least day~2000. These blueshifts, which first appeared around day~600, probably indicate that very dense concentrations of dust persist in the ejecta, although an alternative explanation of asymmetry in the excitation conditions is not ruled out.\\\\ ", + "introduction": "The exceptionally close proximity of the type~IIpec supernova SN~1987A in the Large Magellanic Cloud has provided a unique opportunity to observe a core-collapse supernova with all the resources of modern observational astronomy and over a very long period of time. Near-infrared (near-IR) spectroscopy has played a vital role in the determination of the physical conditions in the debris and the investigation of the element synthesis, providing a valuable complement to optical spectra. In particular, in the nebular phase (when the lines are mostly optically thin) the line profiles produced in the high velocity, homologous expansion have enabled us to examine the ejecta abundances and their {\\it spatial} distribution. During the first few years, all the major southern observatories obtained near-IR spectra of SN~1987A. Observations at the Anglo-Australian Telescope (AAT) covering the first 3~years post-explosion were described in Meikle et al. (1989, 1993) (hereinafter Papers~I, II). \\\\ \\noindent In this paper we describe spectra obtained at the AAT for a further 8 epochs from 1990 November 2 (day 1348 = 3.7~years) to 1995 October 17 (day 3158 = 8.6~years). Epochs are with respect to the explosion date 1987 February 23. Our day~1348 spectrum was the last observation of SN~1987A obtained with the near-IR spectrometer FIGS (Bailey et al. 1988. Also see Papers~I \\& II). All the subsequent spectra were obtained with the more sensitive spectrograph, IRIS (Allen et al. 1993), which allowed the acquisition of high resolution near-IR spectra to a later phase than was achieved at any other observatory. Preliminary reports of these data have been given in Fassia (1999) and Meikle (2001). The only other post-3~year IR spectra reported are those of Bautista et al. (1995) which reached day~1445, but at a lower resolution than those described here. In a future paper (Fassia et al., in preparation) we shall present $HK$-band spectra of SN~1987A taken with the MPE imaging spectrograph, 3D (Weitzel et al. 1996), at the AAT on 1997 Dec 16 (day~3949 = 10.8~years) and 1998 Dec 1 (day~4299 = 11.8~years). We also note that the earliest Hubble Space Telescope (HST) NICMOS spectrum of SN~1987A was taken on 1998 June 15 (day~4130 = 11.3~years), but has not yet been published. ", + "conclusions": "\\subsection{The Physical Framework} The CSM spectra will be discussed elsewhere. Here we confine our discussion to the ejecta spectra. The framework for this section is based upon theoretical descriptions of the physical conditions in the ejecta and line strength predictions at very late times provided by Kozma \\& Fransson (1992), Fransson \\& Kozma (1993), Fransson, Houck \\& Kozma (1996) (FHK96), C97, dKLM98, and KF98a,b. While dKLM98 evolve their model to only 1200~days ({\\it i.e.} 150~days prior to the beginning of our observations) their results are still relevant to the earlier part of the era considered here. KF98a follow the temperature and ionisation evolution of their model to 2000~days, covering a significant portion of our era, but their line emission calculations (KF98b) stop at 1200~days. C97 consider the much later epoch of day~2875, which is near the end of our era. \\\\ \\noindent During the era of the observations described here, the SN~1987A ejecta spectrum is produced in a rather exotic manner. The energy source of the nebula is the radioactive decay of a number of species created in the explosion (Woosley, Pinto \\& Hartmann 1989; Timmes et al. (1996)). These include $^{56}$Co, $^{57}$Co, $^{60}$Co, $^{44}$Ti and $^{22}$Na. All emit $\\gamma$-rays. In addition, $^{56}$Co, $^{44}$Ti and $^{22}$Na emit energetic positrons and $^{60}$Co emits energetic electrons. By day~1348, the most significant energy sources were $^{57}$Co and $^{44}$Ti, contributing about one half and one quarter respectively of the deposited energy (Li, McCray \\& Sunyaev 1993; Timmes et al. 1996). Around day~1565, the dominant source of deposited energy became $^{44}$Ti and by about day~2000 more than half the deposited $^{44}$Ti energy was in the form of positrons. It is generally assumed that positrons are deposited ``on the spot''. Consequently, given the $\\sim60$~year lifetime of $^{44}$Ti, this means that once positron deposition is dominant, the rate of energy deposition approaches a constant value of $1.3\\times10^{36}$ erg/s per $10^{-4}$~M$_\\odot$ of synthesised $^{44}$Ti (Woosley et al. 1989). \\\\ \\noindent In order to reproduce the observed spectra and evolution of SN~1987A at late times, C97, dKLM98 and KF98a,b begin with model nebulae having a number of zones of differing chemical composition. These zones are arranged so as to reflect the strong, deep macroscopic mixing in SN~1987A for which there is a great deal of evidence ({\\it e.g.} Arnett 1988, Pinto \\& Woosley 1988, Shigeyama, Nomoto \\& Hashimoto 1988, Woosley 1988, Haas et al. 1990, Spyromilio, Meikle \\& Allen 1990, Fassia \\& Meikle 1999). Light elements (H, He) are mixed down to low velocities, and iron group elements are mixed up to high velocities. C97, dKLM98 and KF98a,b introduce this mixing in different ways. However, they all invoke a macroscopically mixed ``core'' lying within $\\sim$2000~km/s. (dKLM98 refer to the inwardly mixed H/He zones as the ``inner envelope''.) The core thus contains zones that are H-rich, He-rich, intermediate-element-rich, and iron-group-rich, with the nebula being bathed in the radioactive decay energy. The core is generally of mass 4--6~M$_\\odot$. In addition, C97 and KF98a,b include an outer H-envelope of mass 10~M$_\\odot$ extending out to a velocity of 6000--7000~km/s. \\\\ \\noindent In general, the fraction of the radioactive decay energy that does not directly escape from the nebula is injected into the nebular material via Coulomb interaction with the positrons and Compton-scattered electrons. The energy of the resulting non-thermal high energy electrons then goes towards excitation, ionisation and heating of the nebula (KF98a). Three effects play important roles in the evolution of the temperature and ionisation at the very late times considered here. These are the ionisation/thermal ``freeze-out'' effect, adiabatic cooling, and the ``IR-catastrophe''. The first two are important in the H/He envelope, while the third plays an important role in the metal-rich core. \\subsubsection{Freeze-out} The ionisation freeze-out effect was suggested originally by Clayton et al. (1992) and by Fransson \\& Kozma (1993). These authors pointed out that as the SN evolves, there will eventually come a time when the recombination timescale exceeds the radioactive or expansion timescale, so that the rate of change in the level of ionisation slows significantly. Once this phase is reached, the bolometric luminosity exceeds that of the instantaneous radioactive decay deposition, since some of the luminosity results from recombination following ionisation at a significantly earlier epoch. In the model of KF98a,b the ionisation freeze-out phase begins at 800--900~days in most zones, but has a more pronounced effect on the luminosity in the H/He envelope. However, in their ``inner-envelope'' model, dKLM98 find that the freeze-out does not occur until much later than this. Instead, they identify a ``thermal freeze-out'' where the radiative cooling timescale exceeds that of the expansion timescale. In their model, this occurs before the ionisation freeze-out but, as with ionisation freeze-out, results in a luminosity exceeding that of the instantaneous radioactive decay deposition. There is general agreement that by day~1200, the ejecta are predominantly neutral, but with a fraction of singly-ionised species resulting from the ionisation freeze-out and/or direct ionisation by non-thermal electrons. During 1350--2000~days the electron fraction is $\\sim10^{-3}$ in the H-envelope, rising to 0.1 in the Fe-He zone (KF98a). At 3425~days, Lundqvist et al. (2001) calculate that the fraction of iron that is singly-ionised, is in the range 0.2--0.4. \\subsubsection{Adiabatic cooling} The adiabatic cooling becomes significant when the radiative cooling timescale becomes longer than the expansion timescale. For pure adiabatic cooling, T$\\propto t^{-2}$. Owing to their lower density and metallicity, adiabatic cooling is already dominant as early as $\\sim$250~days in the H-envelope and $\\sim$800~days in the He-envelope (KF98a). For the H/He zones within the core, the higher densities mean that adiabatic cooling begins to dominate at 800-1000~days. While fine-structure line cooling ({\\it cf.} IR catastrophe below) becomes the dominant radiative cooling mechanism in these zones, it never supersedes adiabatic cooling. Nevertheless, fine-structure line cooling may be significant in these regions (C97). KF98a find that the H-envelope temperature lies in the range 400-1000~K on day~1350, falling to a range of 150--300~K by day 2000. The temperatures of the H/He zones within the core fall from $\\sim$900~K to $\\sim$300~K over the same period. The day~2000 KF98a value compares well with the $\\sim$300~K derived from the Balmer continuum by Wang et al. (1996) for about the same time. The model of C97 yields $\\sim$130~K for the H-envelope at 2875~days, compared with 350$\\pm$100~K derived from the Balmer continuum at that time. \\subsubsection{The IR catastrophe} Once the nebular heating/cooling rate drops below a certain level, cooling via low-lying fine-structure transitions dramatically overtakes optical and near-IR transitions as the dominant radiative cooling mechanism. Consequently the stabilising temperature falls abruptly from $\\sim$2000~K to a few $\\times100$~K, and the bulk of the nebula's luminosity shifts to far-IR emission. This effect is known as the ``IR Catastrophe'', and was predicted by Axelrod (1980) in his pioneering work on type~Ia spectral models. Fransson \\& Chevalier (1989) predicted that it could also occur in core-collapse SNe such as SN~1987A. The first direct evidence for this phenomenon occurring in SN~1987A was obtained through the detailed study of the evolution of the near-IR/optical [Fe~II] lines during the second year (Spyromilio \\& Graham 1992). \\\\ \\noindent In the dense, metal-rich zones, fine-structure radiative cooling dominates for a considerable time. [Fe~II]~26~$\\mu$m emission is particularly important. Fine-structure cooling is most pronounced in the Fe-rich zones (KF98a), where it commences at $\\sim$500~days at 2700~K. According to KF98a, by the beginning of our observations (1348~days) the temperature in the Fe-He zone had fallen to $\\sim150$~K, where it remained until 2000~days, after which adiabatic cooling became important. The IR catastrophe is also important for the intermediate mass element zones, where the temperature at 1350~days is 200--400~K, falling to 100--200~K by day~2000 (KF98). Both FK98b and dKLM98 find that, owing to the IR-catastrophe, [Fe~II] 0.72, 1.26, 1.53, 1.64~$\\mu$m lines originating from the newly-synthesised iron in the core should have vanished by $\\sim$600~days. However, detectable flux in these lines persists beyond this epoch. FK98b and dKLM98 attribute this to thermally-excited emission from primordial iron in the H-He envelope where the temperature remains above $\\sim$2000~K until after day~1000. However, as indicated above, by 1350~days, the temperature in the H-He zone is below 1000~K and so thermally excited emission from these lines should be undetectable. \\\\ \\subsection{Near-IR ejecta lines at very late times} We can use the observed near-IR spectra to examine the physical scenario described above. The evolution of the intensities of the principal ejecta lines is illustrated in Figure~8. \\begin{figure*} \\vspace{13cm} \\special{psfile=87a_int.ps angle=0 voffset=-100 hoffset=30 vscale=67 hscale=70} \\caption{Evolution of the line intensities of the most prominent ejecta lines. The boxed point with the error bars indicates the typical $\\pm$35\\% fluxing precision of the ejecta spectra. However, where the spectral model matching is poor, the error in intensity is consequently greater (see text).} \\label{fig_evol} \\end{figure*} \\subsubsection{Hydrogen Lines: Pa$\\beta$, Pa$\\gamma$, Br$\\gamma$} The intensity and evolution of the H-lines in the adiabatic cooling phase is calculated by dKLM98 and KF98b up to 1200~days. For epochs $>2$~years, non-thermal ionisation from the ground state by Compton-scattered electrons dominates photoionisation from excited states (C97, KF98b). Non-thermal excitations are also important at epochs beyond 1000~days (KF98b). dKLM98 find that by day~800 the line ratios should be Case~B. FK98b reach a similar conclusion, but with the transition to Case~B in the core being a little later, at about 1000~days, being complete by day~1200. Thus, during our era, the H-spectrum is predicted to consist of a mixture of a low-temperature Case~B H-recombination spectrum together with a contribution from direct non-thermal excitation (see KF98b for details). Line light curves are given by dKLM98 and KF98b for the near-IR lines Pa$\\alpha$,$\\beta$, Br$\\alpha$,$\\beta$,$\\gamma$,$15$. dKLM98 find that, for the well-observed near-IR H-lines Pa$\\beta$ and Br$\\gamma$, their models makes satisfactory predictions to 1200~days. This is in spite of their non-inclusion of the envelope with its delayed recombination effects. The KF98b model light curve for Br$\\gamma$, which includes delayed recombination effects in the envelope, is also in good agreement with the observations. Their model indicates that, by day~1200, the H-envelope is responsible for $\\sim$70\\% of the Br$\\gamma$ flux, and that this fraction is growing. Both Fransson et al. (1996) and dKLM98 point out that derivation of total H-mass from the H-lines is difficult and subject to large uncertainties. Fransson et al. (1996) and KF98b make use of the line profiles in order to find the most appropriate model (and hence the mass of H). (In a sense, C97 also make use of the observed velocity distribution in determining the H-mass).\\\\ \\noindent We examined the ejecta H-line fluxes to see if they yield the predicted low-temperature Case~B ratios. These were first de-reddened assuming $A_{V}=0.6$ (Blanco et al. 1987). We find that the observed, de-reddened intensity ratio $I_{Pa\\gamma}/I_{Pa\\beta}=0.45-0.6$ for days~1469, 1734, 1822. Martin (1988) predicts a value of $\\sim$0.5 for a temperature of 200--1000~K. Given the uncertainty in the relative line intensities, we judge that the observed ratios are consistent with the predicted values. The uncertainties in the model matching are too large to make a useful judgement about the ratio on epoch day~2112. Martin (1988) also predicts $I_{Br\\gamma}/I_{Pa\\beta}$=0.195. This is in fair agreement with the observed, de-reddened ratio of 0.18 on day~1469. Apparently poorer agreement is obtained on subsequent days. However, this is probably again due to model matching uncertainties, especially in the $K$-window. We conclude that, at least up to $\\sim$1800~days, the observed H-spectra are consistent with the scenario of non-thermal ionisation/excitation within an environment dominated by adiabatic cooling. \\subsubsection{Helium Lines: He~I 1.083, 2.058~$\\mu$m} As for hydrogen, theoretical considerations indicate that post-1000 day He-lines evolve in an environment where the temperature is dominated by adiabatic cooling, and the lines are driven by (a) non-thermal direct excitation, and (b) recombination following non-thermal ionisation. dKLM98 and KF98b point out that the 2.058~$\\mu$m line is uncontaminated by other ejecta species throughout the observations, and also find that it is optically thin in all regions after 700~days. Moreover, after $\\sim$700~days, the 2.058~$\\mu$m line intensity is relatively insensitive to assumptions made about the He~I 584~\\AA\\ continuum destruction probability (KF98b). These considerations lead dKLM98 \\& KF98b to suggest that the post-600~day He~I 2.058~$\\mu$m line can provide the best measure of the total helium mass. However, consideration of our spectra leads us to a more pessimistic view of the usefulness of the 2.058~$\\mu$m line in this situation. By day~1533, it was quite faint. To make matters worse it is blended with comparably-strong CSM flux in this line. Moreover, He~I 2.058~$\\mu$m lies in a bad part of the atmospheric window. Consequently there is large uncertainty in the He~I 2.058~$\\mu$m line flux. dKLM98 show that their model 2.058~$\\mu$m light curve provides a good match to the observations up to 1100~days. The KF98b light curve match is somewhat poorer, with a flux overproduction of about 40\\% between days 800 \\& 1100. The 2.058~$\\mu$m light curves of both dKLM98 and KF98b stop at 1200~days, but suggest a gradual slowing down. For the later epochs described here, to within the uncertainties, the observed fluxes fall on a plausible extrapolation of the dKLM98 model light curve, but continues to fall below that of the (extrapolated) KF98b light curve. \\\\ \\noindent dKLM98 \\& KF98b agree that the He~I 1.083~$\\mu$m line is more difficult to use for the determination of helium abundance. It is more temperature sensitive, and is optically thick for much longer than is the case for He~I 2.058~$\\mu$m. Indeed, KF98b find the 1.083~$\\mu$m line to be optically thick in the He~I region even beyond 2000~days. However, high optical depth in the inner zones may not be too important since, after 1200~days, the 1.083~$\\mu$m emission from the H-zone is predicted to dominate (and was already optically thin at 700 days). KF98b also point out that there may be contamination due to [S~I]~1.082~$\\mu$m. Our empirical model matches also indicate that [S~I]~1.082~$\\mu$m emission may be present. However, this line is part of a multiplet, with another component lying at 1.13~$\\mu$m and having about 0.3 of the intensity. Examination of this part of the spectrum (in spite of it being in a bad part of the atmospheric window) indicates that [S~I]~1.082~$\\mu$m makes, at most, a relatively minor contribution to the 1.08~$\\mu$m ejecta emission (see Fig.~3). We find that the observed fluxes of the He~I 1.083~$\\mu$m line fall on plausible extrapolations of both the dKLM98 and KF98b model light curves. \\subsubsection{Forbidden lines of neutral species: [Si~I], [S~I] [Fe~I]} The low temperature (T$<$400~K) of the electron gas in the macroscopically-mixed core during the period of our observations implies that thermally-excited near-IR ejecta lines of [Si~I], [S~I] \\& [Fe~I] should have faded below detectability. Yet, the [Si~I] and [Fe~I] lines are quite clearly present during this very late phase. (Strong blending makes the presence of the [S~I] lines more ambiguous.) In particular, the [Si~I] 1.645~$\\mu$m and [Fe~I] 1.443~$\\mu$m lines are visible to beyond day~2000. The persistence of these lines provides valuable support for the proposition that, by this era, these lines are produced either by recombination or through direct excitation by non-thermal electrons. \\\\ \\noindent The [Si~I]~1.6~$\\mu$m multiplet is of particular interest here. FHK96 point out that for this species, recombination to the neutral state does not produce any significant line emission in the optical or near-IR region. Consequently, the [Si~I]~1.6~$\\mu$m emission must be due entirely to direct excitation by non-thermal electrons. An interesting consequence of this is that the line luminosity will follow the instantaneous energy input, and will be independent of the ejecta temperature. At very late times this is dominated by the 100\\% absorption of $^{44}$Ti decay positrons. Thus, if the $^{44}$Ti scenario is correct, the luminosity in these lines should converge to a near-constant value. In Figure~8 we see that the [Si~I] 1.645~$\\mu$m line fades by about a factor of 3 between days~1734/1822 and 2112. During this period, the radioactive energy deposition is dominated by $^{44}$Ti and would fade by a factor of about 1.5. Thus, to within the uncertainty in the line intensities, the evolution of [Si~I] 1.645~$\\mu$m is consistent with the radioactive decay and energy deposition being dominated by $^{44}$Ti decay at these late times. We note that FHK96 also show that lines which are driven purely by non-thermal excitation, such as the [Si~I]~1.6~$\\mu$m multiplet, have the potential to provide reliable, temperature-insensitive mass estimates, provided the line profile and the bolometric luminosity is also known. \\\\ \\noindent C97 suggest that, by day~2875, virtually the entire positron luminosity of the $^{44}$Ti is deposited in the Fe/Si-rich clumps. While cooling is mostly by ground-term fine-structure lines, about 10\\% (10$^{35}$~erg/s) is via the UV-optical-NIR lines of neutral species. About 20\\% (of the $10^{35}$ erg/s) is emitted as identified UV lines of Fe~I and Si~I with a further $\\sim$70\\% being down-converted to numerous allowed and forbidden optical/NIR metal lines, forming a quasi-continuum. Of the remaining $\\sim10^{34}$ erg/s, a substantial fraction flows into the [Si~I] 1.6~$\\mu$m and [Fe~I] 1.44~$\\mu$m near-IR multiplets. This appears to be supported by our low-resolution spectra (Figure~6) taken on days~2952 and 3158 where the [Fe~I] 1.443~$\\mu$m and possibly the [Si~I] 1.645~$\\mu$m features are still detected. The luminosity in just the [Fe~I] 1.443~$\\mu$m line on day~2952 is about $0.7\\times10^{34}$ erg/s. \\subsubsection{Forbidden Lines of singly-ionised iron} As with the neutral forbidden lines, the low temperature of the ejecta during the era considered here means that there would have been negligible thermal excitation of the [Fe~II]~1.26, 1.64~$\\mu$m multiplets. This includes the excitation of primordial iron in the H-envelope since, as mentioned above, even there it is expected that the temperature would be less than 1000~K by day~1350 (KF98a). This suggests, therefore, that the persistence of the [Fe~II] lines must be due to recombination or direct excitation by non-thermal electrons. However, we note that C97 state that near-IR [Fe~II] lines produced by radiative cascade are expected to be weak since the major radiative cascade to the ground level goes through optical forbidden lines and FIR lines of the ground term. Detailed modelling is required to test if this is in conflict with our detection of these lines. \\subsection{Velocity behaviour in the ejecta lines} The FWHM velocities of the more prominent ejecta lines are listed in Table~2, Col.~4. The He~I lines apparently exhibit the largest width, at $\\sim$5000~km/s (FWHM). However, the uncertainty in this measurement is large due to strong blending of the He~I ejecta lines with CSM lines and/or other species. The H~I lines and the [Fe~I], [Fe~II] lines show widths of 4000--4500~km/s, while the [Si~I] lines have widths closer to 3000~km/s. Even this latter value is higher than the $\\sim$2000~km/s invoked for the macroscopically mixed core. Of particular interest is the fact that the [Fe~I] and [Fe~II] lines exhibit higher velocity widths than the [Si~I] lines. We argued above that the persistence of the forbidden iron line emission during the era studied here must be driven by delayed recombination following freeze-out or through direct excitation by non-thermal electrons. Consequently, the high velocities in these lines imply that these processes must be occurring well out into the H/He envelope. Detailed modelling will be required to determine how much each process contributes to the line emission. If we favour the latter scenario, it immediately suggests that upward mixing was even greater than has been assumed hitherto. We note that velocities of at least $\\sim$3000~km/s were observed in the [Fe~II] lines as early as the end of year~1 (Paper~I and Spyromilio et al. 1990, Haas {\\it et al.} 1990). Moreover, Fassia \\& Meikle (1999) showed that the presence of the He~I 1.083~$\\mu$m line on days~76--135 implied that upward mixing of $^{56}$Ni in the ejecta of SN~1987A had extended to velocities as high as 3000--4000~km/s. Upward mixing of $^{56}$Ni to even higher velocities (over 5000~km/s) has been recently deduced by Mitchell et al. (2001) on the basis of the high strength of the Balmer lines a few days after explosion. Such high $^{56}$Ni velocities may be evidence of neutrino-instability-driven acceleration of radioactive nickel just after the core-bounce ({\\it e.g.} Herant, Benz \\& Colgate 1992). Alternatively, these results may favour the jet-like explosion models of Nagataki (2000), in which the high velocity [Fe~II] line profiles are well reproduced. \\\\ \\noindent Another remarkable characteristic of these late-time spectra is the presence and persistence of blueshifts (with respect to the supernova rest frame) in the ejecta lines (Table~2, Col.~3), with values of typically --200~km/s to --800~km/s. Such shifts first appeared around day~600 (Paper~II) and were attributed to the formation of dust, blocking out the red (far-side) wing of the line. With the persistence of the blueshifts to as late as day~2000 we conclude that very dense dust concentrations must have formed in the ejecta. However it is possible that asymmetry in the excitation conditions may also be contributing to the effect." + }, + "0112/astro-ph0112502_arXiv.txt": { + "abstract": "Surveys of galaxy clusters provide a promising method of testing models of structure formation in the universe. Within the context of our standard structure formation scenario, surveys provide measurements of the geometry of the universe and the nature of the dark energy and dark matter. Cluster catalogues will be constructed using some combination of X--ray, optical/near--IR, and mm or cm-wave observations. These catalogues will be used to study the cluster redshift and mass distributions along with the correlations of the cluster spatial distribution. These measurements probe the volume--redshift relation, the power spectrum of density fluctuations and the evolution of galaxy cluster abundance. All are sensitive to the amount of dark matter $\\Omega_{M}$, the amount of dark energy $\\Omega_{E}$, the equation of state of the dark energy $w(z)$ and any other parameter, which affects the expansion history of the universe. ", + "introduction": "\\label{sec:intro} Over the last few years, cosmological constraints from Type Ia SNe (Schmidt et al. 1998; Perlmutter et al. 1999), cluster baryon fractions (White et al. 1993a; David et al. 1995; White \\& Fabian 1995; Burles \\& Tytler 1998; Mohr et al. 1999; Arnaud \\& Evrard 1999), the cosmic microwave background (CMB) anisotropy (Hanany et al. 2000; Jaffe et al 2000; Lange et al. 2001) and other complementary measures (Bahcall et al. 1999, and references therein) have pointed toward a dark energy dominated universe ($\\Omega_\\Lambda\\sim{2\\over3}$), with a significant dark matter component ($\\Omega_m\\sim{1\\over3}$) and a trace of baryonic matter. The recent detections of the 2nd and 3rd acoustic peaks in the CMB anisotropy (Halverson et al. 2001; Netterfield et al. 2001; Pryke et al. 2001) lend additional support to these conclusions and bring several important questions into sharp focus. At the dawn of this new era of precision cosmology, the important questions concern the very nature of the dark matter (collisionless or self-interacting) and the characteristics of the dark energy (which we can parametrize by the equation of state parameter $w$, where the pressure $p=w\\rho$). Recent theoretical and experimental developments make future cosmological studies that utilize galaxy clusters extremely promising. One particularly promising approach is the use of galaxy cluster surveys, which enable one to measure the cluster redshift distribution and the correlations in the cluster spatial distribution. Surveys are now being carried out using cluster X--ray emission, the near-IR/optical light from cluster galaxies, the distorted morphologies and alignment of background galaxies, and the effect that hot electrons within clusters have on the cosmic microwave background (the so-called Sunyaev-Zel'dovich effect or SZE; Sunyaev \\& Zel'dovich 1972). To use these surveys to full effect in cosmology studies, we must first test the standard model of structure formation. In addition, we must sharpen our understanding of the nature and evolution of galaxy cluster internal structure and the relationships between cluster observables (i.e. SZE decrement, X-ray emission, galaxy light) and the cluster halo mass. In these proceedings we describe a fundamental test of the hierarchical structure formation model, and then we examine in some detail the cosmological dependences of the cluster redshift distribution. We end by highlighting some of the challenges that currently exist in using cluster surveys to precisely constrain cosmological quantities like the equation of state of the dark energy. ", + "conclusions": "This contribution contains a description of two ways of using cluster surveys to learn about structure formation and cosmology: (1) the context free test of hierarchichal structure formation using SZE cluster surveys, which are sufficiently sensitive to detect low mass clusters no matter what their redshift, and (2) the use of cluster redshift distributions within the context of our standard model for structure formation to determine the quantity and nature of dark matter and dark energy in the universe. It's important to emphasize that there is additional information that comes with a cluster survey. This information allows one to study the cluster mass function $dn/dM$ as a function of redshift, likely improving the constraints derived from integrals over the mass function (i.e. equation~3). In addition, surveys (perhaps with some targeted followup) enable one to study cluster scaling relations such as the X-ray, optical or SZE luminosity--temperature or luminosity--mass relations; a combined study of scaling relations and the redshift distribution may well allow one to solve for the scaling relation evolution and cosmological parameters simultaneously (Diego et al 2001). One can also study the spatial correlations among clusters to infer properties of the underlying power spectrum of dark matter density fluctuations. With good halo mass estimates like those required to use the cluster redshift distribution to full effect, it should be possible to use the cluster power spectrum constraints to improve limits on the neutrino mass density. Even in the absence of accurate halo mass estimates, it should be possible to use large surveys (in volume and number) to measure the scale of the break in the transfer function for the evolution of density perturbations. Recently, Cooray et al (2001) have emphasized that the physical scale of the break in the transfer function, which is the horizon scale at matter--radiation equality, depends on the matter density and CMB temperature. The matter density is measured to high precision with CMB anisotropy observations such as those with MAP and Planck. Therefore, the break in the transfer function is a standard rod, whose scale is independent of redshift and is calibrated to high accuracy with CMB data. Thus, measurements of the cluster correlation function within redshift shells returns the angular diameter distance as a function of redshift, much like the SNe Ia but with a strong physical basis for the lack of evolution in the standard rod. This approach is very complementary to the cluster redshift distribution approach, and it hinges less on extracting unbiased estimates of cluster masses from cluster observables like the X-ray or SZE luminosity." + }, + "0112/astro-ph0112028_arXiv.txt": { + "abstract": "{The oxygen and nitrogen abundances in the H\\,{\\sc ii} regions of the nine Virgo spirals of the sample from \\cite{ski96} and in nine field spiral galaxies are re-determined with the recently suggested P -- method. We confirm that there is an abundance segregation in the sample of Virgo spirals in the sense that the H\\,{\\sc i} deficient Virgo spirals near the core of the cluster have higher oxygen abundances in comparison to the spirals at the periphery of the Virgo cluster. At the same time both the Virgo periphery and core spirals have counterparts among field spirals. Some field spirals have H\\,{\\sc i} to optical radius ratios, similar to that in H\\,{\\sc i} deficient Virgo core spirals. We conclude that if there is a difference in the abundance properties of the Virgo and field spirals, this difference appears to be small and masked by the observational errors. ", + "introduction": "There is evidence that the environment affects the properties of galaxies in clusters (see review of \\cite{bal92}). The most obvious effect is probably that concerning the H\\,{\\sc i} content of galaxies in clusters (\\cite{hay86,huc89,cay94}). It is well established that the spiral galaxies of the Virgo cluster have a tendency to be H\\,{\\sc i} deficient in comparison with normal field spirals (Solanes et al. 1996, 2001) and this deficiency is correlated with distance to the cluster center, the proportion of gas-poor spirals increasing continuously towards the cluster center. Thus, a cluster galaxy (especially a galaxy near the center of the cluster) evolves in a different surrounding gaseous medium in comparison to a field galaxy. The gas exchange between a galaxy and its ambient medium (loss of gas by the galaxy or gas infall onto the galaxy) changes the course of the chemical evolution of the galaxy. Then, it can be expected that the environment affects the chemical properties of galaxies in clusters. Oxygen plays a key role in understanding the (chemical) evolution of galaxies. The origin of oxygen seems to be reliably established in contrast to other elements like carbon or nitrogen. The oxygen abundance can be considered as a tool to investigate the evolution of galaxies. For example, the value of oxygen abundance in a galaxy combined with the value of the gas mass fraction can tell us about the efficiency of mass exchange between a galaxy and its environment (Pilyugin \\& Ferrini 1998, 2000). A number of works have been devoted to searching for the possible effects of cluster environment on the chemical properties of spiral galaxies (\\cite{shi91}; Henry et al. 1992, 1994, 1996; \\cite{ski96}). The conclusions of these works are: {\\it i)} the spirals at the periphery of the cluster are indistinguishable from the field galaxies, {\\it ii)} the H\\,{\\sc i} deficient Virgo core galaxies have larger oxygen abundances at a predetermined galactocentric distance $r = 0.4R_{25}$ (where R$_{25}$ is the isophotal radius) than the field galaxies of comparable luminosity or Hubble type. Accurate oxygen abundances can be derived from measurement of temperature-sensitive line ratios, such as [OIII]4959,5007/[OIII]4363. This method will be referred to as the T$_{e}$ - method. Unfortunately, in oxygen-rich H\\,{\\sc ii} regions the temperature-sensitive lines such as [OIII]4363 are too weak to be detected. For such H\\,{\\sc ii} regions, empirical abundance indicators based on more readily observable lines were some years ago suggested (\\cite{pag79,allo79}). The empirical oxygen abundance indicator R$_{23}$ = ([OII]3727,3729 + [OIII]4959,5007)/H$_{\\beta}$, suggested by \\cite{pag79}, has found widespread acceptance and it has been used for the oxygen abundance determination in H\\,{\\sc ii} regions where the temperature-sensitive lines are undetectable. This method will be referred to as the R$_{23}$ - method. Using the R$_{23}$ - method, the characteristic oxygen abundances (the oxygen abundance at a predetermined galactocentric distance) and radial oxygen abundance gradients were obtained for a large sample of field spiral galaxies (\\cite{vil92,zar94,vzee98}, among others). However, the basic problem whether R$_{23}$ is an accurate abundance indicator is open to discussion (\\cite{zar92,kin94}, among others). It has been found (\\cite{pil00}) that the error in the oxygen abundance derived with the R$_{23}$ -- method involves two parts: the first is a random error and the second is a systematic error depending on the excitation parameter. A new way of oxygen abundance determination in H\\,{\\sc ii} regions (P -- method) has been recently suggested (Pilyugin 2000, 2001a). By comparing oxygen abundances derived through the $T_{e}$ -- method in high-metallicity H\\,{\\sc ii} regions with those derived through the P -- method, it has been found that the precision of oxygen abundance determination with the P -- method is comparable to that of the $T_{e}$ -- method (Pilyugin 2001a,b). It has been also shown that the R$_{23}$ -- method provides more or less realistic oxygen abundances in high-excitation H\\,{\\sc ii} regions but it produces overestimated oxygen abundances in low-excitation H\\,{\\sc ii} regions. Taking into account this fact together with the fact known for a long time (\\cite{sea71,smi75}) that galaxies can show strong radial excitation gradients, in the sense that only the low-excitation H\\,{\\sc ii} regions populate the central parts of some galaxies, one can expect that the oxygen abundances of the inner H\\,{\\sc ii} regions and the gradient slopes based on the (O/H)$_{R_{23}}$ data can be appreciably overestimated. This speculation has been confirmed by comparison of the radial (O/H)$_{R_{23}}$ abundance distribution with the radial (O/H)$_{T_{e}}$ abundance distribution within the disk of the well-observed spiral galaxy M101 (\\cite{pil01b}). The problem whether the cluster environment affects the chemical evolution of galaxies will be considered here based on the (O/H)$_{P}$ and (N/H)$_{P}$ abundances obtained in Section 2. In \\S 3 we analyze these data with the multiphase model applied to three Virgo galaxies, considered as typical examples. A dicussion is included in Section 4 and Conclusions are in Section 5. ", + "conclusions": "The oxygen and nitrogen abundances in H\\,{\\sc ii} regions in the nine Virgo spirals of the sample from \\cite{ski96} and in nine selected field spiral galaxies were re-determined with the recently suggested P -- method. It has been confirmed that there is an abundance segregation in the sample of Virgo spirals in the sense that the H\\,{\\sc i} deficient Virgo spirals near the core of the cluster have higher oxygen abundances in comparison to the spirals at the periphery of the Virgo cluster. In general, both Virgo periphery and Virgo core spirals have counterparts among field spirals. Some field spirals have a small H\\,{\\sc i} to optical radius ratios, similar to the one in H\\,{\\sc i} deficient Virgo core spirals. It has been concluded that if there is an actual difference in abundance properties between the Virgo and field spirals this difference should be small or entirely masked by the errors. These abundance results have been analyzed with the multiphase chemical evolution model, applied to three galaxies NGC~4501, 4321 and 4303, considered as typical examples of galaxies located at the center, at intermediate locations, and at the periphery of the cluster. These models are shown to be able to reproduce the observations of abundances, star formation rate and diffuse and molecular gas densities. Following these models the infall rate was strong at early times and it is very low now for the Virgo core galaxies, but with features continuously smoother for the periphery galaxies. The consequence of the infall of the gas is the dilution of the elements, and therefore lower abundances than those expected with the closed box model, for the gas fractions in the center galaxies, are predicted. Thus, the final abundances turn out to be similar to the ones present in field galaxies. The scenario of the ram pressure stripping usually claimed to produce the observed characteristics in the cluster galaxies, such as a H\\,{\\sc i} deficiency, which is larger for the galaxies closer to the cluster, was recently simulated with N-body models by \\cite{vol01}. Following this work, a large infall rate is produced in the inner regions of the disk of these galaxies, as a consequence of their movement towards the cluster core, with a posterior phase of quenching of the star formation; these results appear consistent with our conclusions." + }, + "0112/astro-ph0112444_arXiv.txt": { + "abstract": "We present the preliminary results of a sub-mm imaging survey of ultracompact HII regions, conducted with the SCUBA bolometer array on the JCMT. ", + "introduction": "Ultracompact (UC) HII regions are currently the best known tracer of massive YSOs and represent the earliest confirmed stage of massive star formation. In excess of 150 UC HII regions have been detected, mainly by radio surveys. Whilst the environments of UC HII regions are known very well on the small scale (a few arcseconds) they are not well known on scales over 40\\arcsec. This is because most UC HIIs have, to date, been observed using either interferometers (to gain information on small scales at the expense of large scales) or by single-position large-beam (typically 40\\arcsec~or worse) spectroscopy. To redress this issue we recently undertook an imaging survey of over 100 UC HII regions using SCUBA on the JCMT, which enables us to rapidly map (with high resolution) the dust emission from the clumps in which the UC HIIs are embedded. ", + "conclusions": "" + }, + "0112/astro-ph0112391_arXiv.txt": { + "abstract": "We review the science case for studying CMB polarization. We then discuss the main issues related to the analysis of forth-coming polarized CMB data, such as those expected from balloon-borne (e.g. BOOMERanG) and satellite (e.g. Planck) experiments. ", + "introduction": "Strong theoretical arguments suggest the presence of fluctuations in the polarized component of the cosmic microwave background (CMB) at a level of 5-10\\% of the temperature anisotropy. A wealth of scientific information is expected to be encoded in this polarized signal. However, while the existence of anisotropies in the temperature of the Cosmic Microwave Background (CMB) has now been firmly established by several experiments \\cite{boom, max, dasi}, only upper limits are currently available for fluctuations in the polarization of the CMB radiation. The prospect of detecting CMB polarization anisotropy at small angular scales is now more promising than in the past. In the next few years, a number of experiments (e.g. BOOMERanG, Planck) will have the right sensitivity, as well as the necessary control on systematic effects, to make the measurement of polarization an achievable goal. In this contribution we will first quickly review the major features of CMB polarization, then we will address some of the issues that will have to be faced in order to analyze the data collected by the forthcoming experiments. ", + "conclusions": "Temperature anisotropy measurements have just started to have the accuracy required for high precision cosmology. Polarization has enormous scientific potential, but is still a big challenge, both experimentally (low signal, fine-scale structure, systematics, etc.) and for data analysis (which must be both accurate and efficient). The next few years will most likely bring us definitive high-resolution high-sensitivity maps of the CMB temperature anisotropy by satellites (MAP, Planck). The new frontier of cosmological exploration will then shift towards observations of CMB polarization, which will certainly provide us with new insights about the physics of the early universe. \\begin{figure}[!ht] \\resizebox{.95\\textwidth}{!} {\\includegraphics{ib.ps} \\includegraphics{i.ps}} \\end{figure} \\begin{figure}[!ht] \\resizebox{.95\\textwidth}{!} {\\includegraphics{qb.ps} \\includegraphics{q.ps}} \\end{figure} \\begin{figure}[!ht] \\resizebox{.95\\textwidth}{!} {\\includegraphics{ub.ps} \\includegraphics{u.ps}} \\caption{Simulated maps for the 100 GHz channels (32 radiometers) of Planck/LFI. Shown from top to bottom are the I, Q and U components, obtained by a naive coadding of observations in each pixel (left) and using the map-making procedure described in the text (right).}\\label{mapmak} \\end{figure} \\begin{theacknowledgments} It is a pleasure to thank the organizers of this interesting workshop for the invitation and for the stimulating environment. We acknowledge use of HEALPix ({\\tt http://www.eso.org/science/healpix/}) and CMBFAST. \\end{theacknowledgments}" + }, + "0112/astro-ph0112358_arXiv.txt": { + "abstract": "I summarize recent work comparing relative distances measured to individual galaxies with independent methods. The comparisons include: ground-based surface brightness fluctuation (SBF) and fundamental plane distances to 170 galaxies, distances predicted from galaxy velocities and the inferred gravity field, HST SBF measurements to seven early-type hosts of Type Ia supernovae, and ties of the Cepheid distance scale to early-type galaxies. Independent calibrations for some methods provide interesting constraints on the Cepheid zero point. ", + "introduction": "The two most frequently applied early-type galaxy distance indicators are the fundamental plane (FP, and the related $D_n$-$\\sigma$) and surface brightness fluctuations (SBF) methods. In a recent study (Blakeslee \\etal\\ 2001,\\,2002), we used $V$- and $I$-band data from the ground-based SBF Survey (Tonry \\etal\\ 2001) to calculate FP photometric parameters for 170 galaxies with velocity dispersions available in the homogenized SMAC catalogue (Hudson \\etal\\ 2001). To our knowledge, this is the largest galaxy-by-galaxy comparison of different standard candle/rod distance methods to date. Fig.\\,1a shows the comparison.~~ Overall the distance agreement was good, but several low-luminosity, S0 galaxies had systematically low FP distances, probably due in part to younger ages and lower mass-to-light ratios, although aperture effects may also contribute. The SBF distances are tied to the Cepheids via measurements in spiral bulges, while the FP distances are tied to the Hubble flow via distant clusters; the Hubble constant that results from this comparison is $H_0=68$ \\kmsM. However, we also derived independent distances for these galaxies based on their velocities and the gravity field inferred from the redshift-space galaxy density; the resulting comparison with SBF yields $H_0=74$ (Fig.\\,1b), formally discordant at the 2$\\,\\sigma$ level with the FP-SBF result, but within the range of the systematic uncertainties in the various ties. Another interesting facet of this work relates to the ``fluctuation number'' $\\Nbar\\equiv \\mbar-m_{\\rm tot}$, which measures the galaxy luminosity in units of the weighted mean stellar luminosity. \\Nbar\\ correlates tightly with stellar velocity dispersion; it also correlates with galaxy color and is independent of Galactic extinction. Interestingly, SBF distances calibrated using the properties of \\Nbar, such as those shown in Fig.\\,1b, can be viewed as a hybrid of SBF and FP distances, and may be more accurate than those calibrated from galaxy color alone. We plan to investigate these issues in more detail. \\begin{figure}[t] \\plottwo{fpsbf68d.ps}{nbariras74d.ps} \\caption{\\footnotesize{\\bf (a)} Comparison of FP and SBF distance moduli for the 170 galaxies in the cross-matched SBF-SMAC survey samples (from Blakeslee \\etal\\ 2002). The lower panel shows distance residuals. Filled circles represent true ellipticals, while open circles represent S0s. Six galaxies having systematically uncertain FP or SBF distances are shown as crosses. {\\bf (b)}~Same as (a), but for the comparison of \\Nbar-calibrated SBF distances with those predicted from the observed galaxy density field in redshift space (Virgo core galaxies have been assigned the systemic velocity). \\vspace{-7pt} } \\end{figure} \\vspace{-3pt} ", + "conclusions": "" + }, + "0112/astro-ph0112322_arXiv.txt": { + "abstract": "X-ray images of some young supernova remnants show bright point sources which have not been detected in radio, optical and gamma-ray bands. Despite the similarity of the X-ray spectra of these objects, they show a variety of temporal properties. Most likely, they are neutron stars whose properties (spin periods? magnetic fields? environments?) are different from those of radio and/or gamma-ray pulsars. We present an overview of observational results on several objects of this class --- the central sources of Cassiopeia A, RX J0852--4622, RCW 103, Puppis A, and PKS 1209--51/52 --- with emphasis on the recent {\\sl Chandra} observations. ", + "introduction": "There are two types of isolated (non-binary) neutron stars (INSs) with quite different observational manifestations. {\\it Active, rotation-powered pulsars} include, in addition to commonly known radio pulsars, some ``radio-quiet'' INSs whose radio pulsations are not seen because of unfavorably directed pulsar beams. Pulsar activity in these radio-quiet active pulsars can be seen at wavelengths other than radio (e.g., in $\\gamma$-rays --- the classical example is Geminga), and it can manifest itself, in young pulsars, through a compact pulsar-wind nebula (PWN). A distinctive feature of active pulsars is that these objects are powerful sources of {\\em nonthermal radiation from NS magnetospheres}, in radio through $\\gamma$-rays. The intensity of this nonthermal radiation usually exceeds that of the surface (thermal) radiation, at least in young pulsars. Recent observations have established that there exist isolated compact objects, presumably INSs, which show no signs of pulsar activity --- i.e., the usual pulsar processes, which result in beams of relativistic particles, PWNe, and strong nonthermal radiation, do not operate in these objects. They are often called {\\it radio-silent INSs} (or INS candidates), {\\it failed pulsars}, {\\it dead pulsars}, etc --- none of these terms describes their properties adequately (e.g., some of these objects do show X-ray pulsations, but their origin is apparently quite different from that observed from `ordinary' pulsars). These objects can be further divided in at least four classes, based on their observational manifestations. First, six of them have been dubbed {\\it Anomalous X-ray Pulsars} (AXPs; Mereghetti \\& Stella 1995; van Paradijs, Taam, \\& van den Heuvel 1995) because they resemble accreting binary pulsars, but their luminosities, $L_{\\rm x}\\sim 10^{35}$ erg s$^{-1}$, are two-three orders of magnitude lower, their spectra are considerably softer than those of binary pulsars, and their periods cluster in a narrow range of 6--12 s (Israel, Mereghetti, \\& Stella 2001). The second class is comprised of three (perhaps, four) {\\it Soft Gamma-ray Repeaters} (SGRs; Kouveliotou 1995) whose quiescent radiation shows many similarities with AXPs, but, contrary to AXPs, SGRs occasionally undergo strong outbursts, with peak luminosities up to $10^{44}$ erg s$^{-1}$. Thompson \\& Duncan (1996) suggested that AXPs and SGRs are {\\it magnetars} --- INSs with superstrong magnetic fields, $\\sim 10^{14}$--$10^{15}$ G. Some of the AXPs and SGRs are apparently associated with supernova remnants (SNRs). Recent results on AXPs and SGRs are presented by S.\\ Kulkarni (this volume). One more class of presumably older ``{\\it truly isolated}'' (i.e., not associated with SNRs), radio-silent NSs has emerged recently (e.g., Treves et al.\\ 2000). These objects show very soft, thermal-like X-ray radiation with temperatures $kT\\sim 50$--150 eV and luminosities $\\sim 10^{32}$--$10^{33}$ erg s$^{-1}$. Finally, {\\it Compact Central Objects} (CCOs) have been found in several SNRs, which have not been identified as active pulsars, AXPs or SGRs. These objects are particularly puzzling. Below we will discuss recent results on five CCOs which have been observed with the {\\sl Chandra} X-ray Observatory. ", + "conclusions": "Although some properties of CCOs (in particular, their spectra) are very similar to each other, this still does not guarantee that they represent a uniform class of objects. The most outstanding among these sources is the RCW 103 CCO, with its highly variable flux (time scales hours to months), putative 6-hr period, and a possible IR counterpart. We presume that this source is not a truly isolated NS (i.e., it is powered by accretion, at least in its high state), although it is not a rotation-powered active pulsar. The other four CCOs have shown neither long-term variabilities nor indications of binarity. If we adopt a plausible hypothesis that their radiation is thermal and assume that it is adequately described by the same spectral model (e.g., BB or H atmosphere) for each of the sources, then we have to conclude that the emitting area is growing with the CCO's age. If we assume that the radiation emerges from a hydrogen atmosphere, than the size is consistent with a NS radius for the two oldest CCOs, being smaller for the two younger ones. One might suggest that the two younger and two older CCOs are objects of different types, but it does not answer the fundamental question: {\\it Why are the emitting regions of at least two young CCOs so small?} If this is due to the above-discussed nonuniformities of magnetic field or chemical composition, one could speculate that these nonuformities are formed at very early stages of the NS life, and they weaken (spread over the NS surface) in a time of several thousand years. One could assume that the two younger CCOs are BHs, not NSs, but then we would have to invoke accretion as an energy source, which does not look consistent with a lack of variability. The key property to understand the nature of CCOs is their periodicity. So far, the period has been found with high significance (albeit in only one observation) only for the oldest of the CCOs, and its value of 424 ms proves that at least this CCO is a NS. The suspected 12 ms period of the youngest CCO is very puzzling --- we have to wait for its confirmation before making definitive conclusions on the nature of this source. Another intriguing question is whether the CCOs are close relatives of AXPs and SGRs, as one could assume based on similarity of their spectra. Based on the above-mentioned CCO periods (which, however, require confirmation), the formal answer should be `no' rather than `yes'. On the other hand, it is hard to explain why the apparent temperatures and sizes of emitting regions would be so similar in objects of different nature. Therefore, we do not rule out the possibility that at least some of the CCOs belong to the same family as AXPs and/or SGRs, perhaps at different stages of their evolution. In this respect, it would be particularly interesting not only to confirm/measure the CCO periods, but also to evaluate the period derivatives. To conclude, with the aid of the {\\sl Chandra} observations two more CCOs have been discovered, data of much better quality were obtained, and the main problems were formulated more clearly. We expect these problems to be resolved by further observations --- X-ray timing and spectroscopy, supplemented with deep NIR imaging." + }, + "0112/astro-ph0112052_arXiv.txt": { + "abstract": "s{Since the discovery of the Cosmic Microwave Background (CMB) in 1965, characterization of the CMB anisotropy angular power spectrum has become somewhat of a holy grail for experimental cosmology. Because CMB anisotropy measurements are difficult, the full potential of the CMB is only now being realized. Improvements in experimental techniques and detector technology have yielded an explosion of progress in the past couple of years resulting in the ability to use measurements of the CMB to place meaningful constraints on cosmological parameters. In this review, I discuss the theory behind the CMB but focus primarily on the experiments, reviewing briefly the history of CMB anisotropy measurements and focusing on the recent experiments that have revolutionized this field. Results from these modern experiments are reviewed and the cosmological implications discussed. I conclude with brief comments about the future of CMB physics.} ", + "introduction": " ", + "conclusions": "" + }, + "0112/astro-ph0112264_arXiv.txt": { + "abstract": "We have carried out further analysis of the tentative, short-term brightenings reported by Sahu et al.\\ (2001), which were suggested to be possible lensings of Galactic-bulge stars by free-floating planets in the globular cluster M22. Closer examination shows that---unlikely as it may seem---small, point-like cosmic rays had hit very close to the same star in both of a pair of cosmic-ray-split images, which cause the apparent brightenings of stars at the times and locations reported. We show that the observed number of double hits is consistent with the frequency of cosmic rays in WFPC2 images, given the number of stars and epochs observed. Finally, we point to ways in which cosmic rays can be more directly distinguished. ", + "introduction": "Sahu et al.\\ (2001) have recently reported observations of microlensing of stars in the Galactic bulge by stars of the globular cluster M22. They report one major event, with a characteristic time of $\\sim$18 days and a brightening by 3 magnitudes, and six brightenings of 0.3 to 0.8 magnitudes, each seen as similar brightenings in both images of a pair of images taken 6 minutes apart. We discuss here only the short-term events. We demonstrate that each pair of brightenings is caused by two separate cosmic ray hits, one in each image of the CR-split, that happen to occur near the same star. In \\S\\ 2 we describe the observations and our re-analysis; in \\S\\ 3 we present the evidence that the short-term events are caused by cosmic rays; in \\S\\ 4 we discuss the prevalence of CRs, and in \\S\\ 5 we suggest effective ways of reliably avoiding CR contamination. Since there are already at least five papers (Gaudi 2001, de la Fuente Marcos \\& de la Fuente Marcos 2001, Hurley \\& Shara 2001, Fregeau et al.\\ 2001, Soker, Rappaport, \\& Fregau 2001) that discuss the short-term events, we feel that it is urgent to report our new finding, based on further analysis. ", + "conclusions": "We conclude that the 6 minor events found in the GO 7615 observations of M22 can best be explained by coincident cosmic rays rather than by gravitational lensing. Indeed, Sahu et al.\\ (2001) stated that ``The interpretation of these events as microlensing is necessarily tentative.'' That caution was even more appropriate than it seemed at the time. Although these apparent brightenings are not caused by microlensing, we should note that microlensing remains a sensitive technique to detect the presence of small-mass objects in a globular cluster." + }, + "0112/astro-ph0112058_arXiv.txt": { + "abstract": " ", + "introduction": "It is generally agreed that the tremendous densities reached in the centers of neutron stars provide a high pressure environment in which numerous particles processes are likely to compete with each other. These processes range from the generation of hyperons to quark deconfinement to the formation of kaon condensates and H-matter \\cite{weber99:book}. Another striking possibility concerns the formation of absolutely stable strange quark matter, a configuration of matter even more stable than the most stable atomic nucleus, iron. In the latter event all neutron stars would in fact be strange (quark matter) stars \\cite{madsen98:b}, objects largely composed of pure strange quark matter, eventually enveloped in a thin nuclear crust made up of ordinary, hadronic matter. There has been much recent progress in our understanding of quark matter, culminating in the discovery that if quark matter exists it will be in a color superconducting state \\cite{alford98:a,rapp98:a,rajagopal01:a,alford01:a}. The phase diagram of such matter appears to be very complex \\cite{rajagopal01:a,alford01:a}. At asymptotic densities the ground state of QCD with a vanishing strange quark mass is the color-flavor locked (CFL) phase. This phase is electrically neutral in bulk for a significant range of chemical potentials and strange quark masses \\cite{rajagopal01:b}. If the strange quark mass is heavy enough to be ignored, then up and down quarks may pair in the two-flavor superconducting (2SC) phase. Other possible condensation patters are the recently discovered CFL--$K^0$ phase \\cite{bedaque01:a} and the color-spin locked (2SC+s) phase \\cite{schaefer00:a}. The magnitude of the gap energy lies between $\\sim 50$ and $100$~MeV. Color superconductivity thus modifies the equation of state (eos) at the order $(\\Delta / \\mu)^2$ level, which is only a few percent. Such small effects can be safely neglected in present determinations of models for the eos of neutron stars and strange quark matter stars. There has been much recent work on how color superconductivity in neutron stars could affect their properties \\cite{rajagopal01:a,alford01:a,rajagopal00:a,alford00:a,alford00:b}. These studies revealed that possible signatures include the cooling by neutrino emission, the pattern of the arrival times of supernova neutrinos, the evolution of neutron star magnetic fields, rotational (r-mode) instabilities, and glitches in rotation frequencies. In this review I shall complement this list by reviewing several, most intruigung astrophysical implications connected with the possible absolute stability of strange quark matter. (Surface properties of strange matter are discussed in Usov's paper elsewhere in this volume.) This is followed by a discussion of two astrophysical signals that may point at the existence of quark matter in both isolated neutron stars as well as in neutron stars in low-mass x-ray binaries (LMXBs). We recall that a convincing discovery of quark matter in neutron stars would demonstrate that strange quark matter is not absolutely stable, ruling out the absolute stability of strange quark matter and the existence of strange quark stars, for it is not possible for neutron stars to contain quark matter cores and strange matte quark stars to both be stable \\cite{rajagopal01:a}. \\goodbreak ", + "conclusions": "" + }, + "0112/astro-ph0112434_arXiv.txt": { + "abstract": "We point out that the tidal disruption of a giant may leave a luminous ($10^{35}-10^{39}$ erg/s), hot ($10-100$ eV) stellar core. The ``supersoft\" source (SSS) detected by {\\it Chandra} at the center of M31 may be such a core; whether or not it is, the observations have shown that such a core is detectable, even in the center of a galaxy. We therefore explore the range of expected observational signatures, and how they may be used to (1) test the hypothesis that the M31 source is a remnant of tidal stripping, and (2) discover evidence of \\bh s and disruption events in other galaxies. ", + "introduction": "A star of mass $M_\\ast$ and radius $R_\\ast$ can be tidally disrupted by approaching within $R_t$ of a black hole (BH). \\begin{equation} R_t=\\Bigg( \\eta^2 {M_{bh} \\over {M_\\ast}} \\Bigg)^{{1}\\over{3}}\\, R_\\ast, \\end{equation} $M_{bh}$ is the BH's mass; $\\eta$ is a parameter of order of unity. Let $R_S$ be the BH's Schwarzschild radius. For $M_{bh}$ \\gax\\ $10^8\\, M_\\odot,$ $R_t < R_{S}$ for most main-sequence stars, and only tidal disruptions (TDs) of giants lead to observable effects. The rate of such TDs, $\\dot N_{G},$ depends on $M_{bh}$, the kinematics of stars near the galactic nucleus, and the mass and age distribution of these stars (see, e.g., Magorrian \\& Tremaine 1999). Computed rates are typically in the range $10^{-6}-10^{-4}$ yr$^{-1}$. A convenient expression is \\begin{equation} \\dot N_{G} \\sim 10^{-5} \\Big({{L}\\over{L_\\ast}}\\Big)^{1.2} yr^{-1}, \\end{equation} with $L_\\ast =1.8 \\times 10^{10} L_\\odot;$ this was derived by Syer \\& Ulmer (1999), who also provide cautions about the use of a simple scaling law. Studies of TD have focused on possible associated accretion events, lasting for months or decades, with luminosities as high as $\\sim 10^{44}-10^{46}$ erg/s (see, e.g., Hills 1975, Lidskii \\& Ozernoi 1979, Gurzadyan \\& Ozernoi 1980, Rees 1988, Loeb \\& Ulmer 1997, Ulmer, Paczy\\'nski, \\& Goodman 1998, Ulmer 1999). \\noindent{\\bf The Hot Core of a Disrupted Giant:\\ } The disruption leaves an end-product, the giant's hot dense core, whose presence, influence, and observability are the subjects of this paper. The core remains hot ($T>10^5$ K) and bright ($L>10^2 L_\\odot$) for $10^3-10^6$ years, thereby providing the longest-lasting signal of a TD. \\noindent{\\bf Detectability:\\ } The soft X-ray sensitivities of {\\it Einstein}, ROSAT, {\\it Chandra}, and XMM-{\\it Newton} have allowed them to detect supersoft X-ray sources (SSSs) in other galaxies. SSSs (Greiner et\\,al. 1991, Rappaport, \\rd\\ \\& Smith 1994, Kahabka \\& van den Heuvel 1997) have luminosities ($10^{35}-10^{38}$ erg/s) and temperatures ($k\\, T \\sim 10-100$ eV), comparable to those of hot stellar cores. {\\it Chandra's} 0.5--1\\asec\\ angular resolution allows individual sources to be resolved, even in the dense central regions of nearby galaxies. \\noindent{\\bf The Example of M31:\\ } We have studied 8.8 ksec of data collected by Chandra's ACIS-I detector. To identify SSSs we applied a hardness ratio test to all point sources in a $16'\\times 16'$ region centered on the center of M31, selecting all sources with more than 50\\% of their photons below 0.7 keV. Three sources satisfied this condition. Remarkably, one of these is coincident with the center of M31. The luminosity is $\\sim 10^{37}$ erg/s; the data are consistent with little or no emission above $1.5$ keV. The source appears to be variable; data from earlier missions ({\\it Einstein}, ROSAT), e.g., are consistent with the source providing zero flux on one occasion. Garcia {\\it et al.} (2000) have argued that the density of point X-ray sources in this field is so low that a chance coincidence with the center is highly unlikely; this argument is even stronger when applied to SSSs. The observed SSS-like behavior is therefore most likely related to the environment of the center, presumably to the presence of a MBH. The data do not seem consistent with ADAF disk models (see, e.g., Garcia {\\it et al.} 2000), but are compatible with the signature of a hot stellar core. Whether or not the source is a hot core, the M31 observations establish that it is possible to detect such an object in the centers of nearby galaxies. ", + "conclusions": "When a giant is tidally disrupted, the stripped core may provide a hot and bright signature for $10^3-10^6$ years. Detecting such cores and relating them to TD events is challenging. Fortunately, the present generation of instruments, including {\\it Chandra,} XMM-{\\it Newton,} and HST, promise opportunities for progress. This new method of studying MBHs complements searches for flare events that may be due to TDs. The latter must of course be used for distant galaxies, and there is a growing body of evidence for luminous UV and X-ray events lasting for months to years, which may be consistent with the accretion of the envelopes of tidally-disrupted stars.~\\footnote{Examples: the UV flare in NGC 4552 (Renzini {\\it et al.} 1995) the $\\sim 5\\times 10^{43}$ erg/s outburst in IC 3599 (Brandt et al. 1995, Grupe et al. 1995), the $\\sim 2\\times 10^{43}$ erg/s outburst in NGC 5905 (Bade et al. 1996, Komossa \\& Bade 1999), the $>\\, 9\\times 10^{43}$ erg/s flare in RXJ\\,1242.6-1119 (Komossa \\& Greiner 1999), the $\\sim 10^{44}$ erg/s flare in RXJ\\,1624.9+7554 (Grupe et al. 1999), the $\\sim 10^{44}$ erg/s flare in RXJ\\,1331.9-3243 (Reiprich \\& Greiner 2000), or the $\\sim 2\\times 10^{44}$ erg/s flare in RX J1420.4+5334 (Greiner et al. 2000).} Such events are so rare, however, that we would be lucky to detect one in the near future among the small number of galaxies within $5$ or $10$ Mpc. In nearby galaxies, however, the long-lived signatures of the stripped cores may allow us to infer that TDs have occurred. The two complementary modes of study should help us to better understand the frequency and consequences of tidal disruption events and to use them as probes of massive black holes in the center of galaxies." + }, + "0112/astro-ph0112352_arXiv.txt": { + "abstract": "In an effort to better understand the UV properties of ultraluminous infrared galaxies (ULIGs), and compare them to the rest-frame UV properties of high redshift sub-mm and Lyman-break galaxies, we have obtained far- and near-UV imaging observations ($\\lambda_{eff}=1457\\mbox{\\AA}$, $\\lambda_{eff}=2364\\mbox{\\AA}$, respectively) of two luminous infrared galaxies (LIGs--VV 114 and IC 883) and five ULIGs (IRAS 08572+3915, Mrk 273, IRAS 15250+3609, Arp 220, and IRAS 19254--7245) using the Hubble Space Telescope. All the galaxies were detected in both channels. UV light, both diffuse and from star clusters, can be traced to within the inner kpc of the dominant near-IR nuclei. However, in general, the brightest UV sources are clearly displaced from the $I$-band and near-IR peaks by at least hundreds of pc. Further, only 0.07\\%-7.3\\% of the total near-UV light is projected within the inner 500 pc radius, even though this is the same region wherein most of the bolometric energy is generated. All nuclei are highly obscured by dust. Even after correction for dust reddening, the global UV emission fails to account for the total bolometric luminosities of these systems by factors of 3--75. The discrepancy is much worse if only the central regions, where the bolometric luminosities are generated, are included. In two cases (VV 114 and IRAS 08572+3915), the merging companion galaxies are more prominent in the UV than the more IR-luminous member. While all our galaxies show possible signatures of AGN activity, only IRAS 19254--7245 yields even a possible detection of an AGN in our UV images. Simple calculations show that all but one of our galaxies would be expected to drop below the detection thresholds of, e.g., the Hubble Deep Fields at redshifts between 1.5 and 3, and we find that $\\sim$2 of our 5 ULIGs would be selected as Extremely Red Objects in this redshift range. A typical ULIG in our sample would be too faint to be detected at high-redshift in the deepest current optical or sub-mm deep surveys. Only VV 114 has UV luminosity and color similar to Lyman-break galaxies at $z\\sim 3$; the other galaxies would be too faint and/or red to be selected by current surveys. The low UV brightnesses of our ULIGs mean that they would not appear as optically-bright (or bright ERO) sub-mm galaxy counterparts, though they might be similar to the fainter sub-mm galaxy counterparts. ", + "introduction": "A significant challenge in astronomy today is understanding the relationships between very distant galaxies and nearby galaxies. Recent observations have identified many luminous, high-redshift systems, yet they are faint and very difficult to study in detail. Are there good local counterparts for these galaxies? If so, we may try and use the local counterparts to better understand their distant relatives. New observations, carried out primarily with the Sub-millimeter Common User Bolometer Array (SCUBA: Holland et al. 1999) on the James Clerk Maxwell Telescope, have revealed the presence of a significant population of far-IR (FIR) luminous galaxies, most likely at redshifts $z \\ga 2$ (``sub-mm galaxies,\" e.g., Smail et al. 1998, Barger et al. 1999a). These galaxies are apparently responsible for much of the FIR and sub-mm extragalactic background radiation, whose energy density is comparable to that of the integrated optical light of faint galaxies seen in, e.g., the Hubble Deep Fields (HDFs) (e.g., Hauser et al. 1998; Barger, Cowie \\& Sanders 1999). The sub-mm galaxies detected to date are much more luminous than almost any galaxy in the local universe. It is possible that a significant fraction of the star formation in the early universe occurred in FIR-luminous galaxies, of which the current sample of sub-mm galaxies are the most luminous. At present, it is not clear if the sub-mm galaxies are at all related to galaxies at similar redshifts discovered by the Lyman break technique, even though the brightest, reddest Lyman break galaxies (LBGs) may have comparable bolometric luminosities, after reddening is taken into account \\citep{meu99,ade00}. Most (but not all) sub-mm galaxies are too faint and/or red to be found with LBG selection criteria. For the sub-mm galaxies with known redshifts, the bluest sub-mm observations (at, e.g., 450 \\micron) still fall on the Rayleigh-Jeans side of the rest-frame FIR peak wavelengths. Nonetheless, the limited information on these sub-mm galaxies shows that they appear to have overall rest-frame FIR/radio spectral energy distributions quite similar to those of the ultraluminous infrared galaxies discovered in the local universe ($z\\la 0.1$). This has led to increased interest in the nearby infrared galaxies as possible counterparts. Ultraluminous infrared galaxies \\footnote{See \\citet{san96} for a review. $L_{IR}=L(8-1000\\micron)$, computed using the prescription in that paper.} (ULIGs, with $L_{IR}>10^{12}L_{\\sun}$) were found to be a significant class of objects upon analysis of the Infrared Astronomical Satellite (IRAS) all-sky survey \\citep{soi87, san88}. The bolometric luminosities of ULIGs, the most powerful infrared-luminous galaxies, are similar to those of optically-selected quasars. We now understand ULIGs to be a stage in the merger of two gas-rich spiral galaxies. During the merger, the molecular gas is driven inward towards the nuclei, where the gas serves to fuel a tremendous burst of star formation and, perhaps, AGN activity as well. Dust absorbs and re-radiates the light from young stars and/or the AGN. Complicating our understanding of ULIGs is the fact that most of the energetic activity takes place in the inner few hundred pc, where small angular sizes and almost entirely opaque ISM prevent us from directly witnessing events occurring in the nuclei themselves. The primary goal of our study is to explore the possible connection between sub-mm galaxies and ULIGs. There is substantial circumstantial evidence linking ULIGs and sub-mm galaxies, but a direct connection has not been established. In part, this is because of observational considerations. The FIR/radio spectral energy distributions (SEDs) of the best-studied sub-mm galaxies (ERO J164502+4626.4: Dey et al. 1999; SMM J02399--0136: Ivison et al. 1998; SMM J14011+0252: Ivison et al. 2000; and SMM J02399--0134: Barger et al. 1999a) are sampled redward of the rest-frame FIR peak, which should occur at 60--100 \\micron\\ if their dust temperatures are similar to those of the ULIGs. Optical and near-IR observations sample blueward of $\\sim$0.6 \\micron\\ in the rest-frame for $z\\sim 2.5$. As a result, we know little or nothing of the properties of sub-mm galaxies from the red optical through $\\sim$200 \\micron\\ in the rest-frame. And unfortunately, local ULIGs are poorly studied shortward of $\\sim$4000 \\AA. What morphological comparisons we can make using, e.g., optical images, are quite limited. As a result, the strongest links between ULIGs and sub-mm galaxies are made largely on the basis of three facts. First, ULIGs and sub-mm galaxies are the most highly luminous systems in the far-IR in the local universe and at $z \\ga 1$ respectively, although most sub-mm galaxies detected so far are inferred to be still a factor of few more luminous than ULIGS such as Arp 220 (Barger, Cowie \\&\\ Richards 2000). Second, ULIGs and sub-mm galaxies also have similar far-IR/sub-mm/radio SEDs. Finally, large amounts of molecular gas (e.g., CO) have been detected in both classes of objects. Local counterparts, even when not exact matches, can offer important insights on high-{\\em z} systems. For example, starbursts observed with IUE (e.g., Meurer et al. 1999 and references therein) have SEDs and UV spectral properties like LBGs, though the local galaxies are smaller and less luminous; this has been used to infer that the basic physics (e.g., winds, reddening) of these objects are similar. We hope that observations of ULIGs in the rest-frame UV, telling us about their morphologies, luminosities, and colors, might both allow a test of their fitness as counterparts of sub-mm galaxies, and also allow us to better interpret the rest-frame UV observations of sub-mm galaxies. A secondary goal is to test whether ULIGs follow what is known as the ``IRX-$\\beta$\" correlation between the redness of the UV continuum (parameterized by the UV spectral slope $\\beta$ between $\\sim1600\\mbox{\\AA}$ and $2200\\mbox{\\AA}$, where $f_\\lambda \\propto \\lambda^{\\beta}$) and the IR-excess (the IRX, defined as the ratio of the FIR/UV fluxes\\footnote{In this paper, we compute the FIR fluxes according to the prescription of Helou et al. 1988, including the dust temperature dependent color correction factor defined in the appendix of that paper. We note that Meurer et al. use a constant color correction factor reasonable for the typical dust temperatures of their galaxies. The UV flux is computed as $F_{\\lambda}=\\lambda \\times f_{\\lambda}$ in the FUV channel, where $f_\\lambda$ is the observed flux density in erg/cm$^2$/s/$\\mbox{\\AA}$.}), that has been established for a sample of local starbursts observed in the UV \\citep{meu95,meu97,meu99}. The correlation is well fit by a foreground dust shell model, wherein dust surrounding the UV-bright region both reddens and absorbs the UV light, reprocessing it to FIR wavelengths. The correlation appears to hold for $L\\la 10^{11.5} L_{\\odot}$, but has not been tested at higher luminosities. If ULIGs also follow this correlation, then we would have more confidence that rest-frame UV observations of even the most dusty high-$z$ star-forming galaxies can be used to estimate their total bolometric luminosities (dominated by the difficult to access far-infrared) ensuring that these galaxies are included in the census of star formation in the early universe. We know that $\\beta$ and the IRX correlate with total luminosity for the Meurer et al. sample \\citep{hec98}. If these trends hold at higher $L$ and higher $z$, then sub-mm galaxies may be too red to be selected as LBGs, and may perhaps better be associated with Extremely Red Objects (e.g.\\ Ivison et al.\\ 2000, ApJ, 542, 27). In this paper, we present new rest-frame UV observations of seven galaxies with $L_{IR}>10^{11.5}$. Three ULIGs (VII Zw 31, IRAS 12112+0305, and IRAS 22491--1808) were detected in the UV using the Faint Object Camera on the pre-COSTAR Hubble Space Telescope (HST) by Trentham, Kormendy, \\& Sanders (1999, hereafter TKS). Though their photometric precision was limited by the spherical aberration and high backgrounds, the detections were good evidence that our program was in fact feasible. More recently, ground-based images in the $U^{\\prime}$ band ($\\lambda=3410\\mbox{\\AA}$, $\\Delta\\lambda=320\\mbox{\\AA}$) were obtained for many ULIGs (including three of our galaxies, IRAS 08572+3915, Mrk 273, and IRAS 15250+3609, plus two galaxies, IRAS 12112+0305 and IRAS 22491--1808, from TKS) by \\citet{sur00a}. This paper concentrates on the large-scale photometric and morphological properties for the galaxies in our sample. We will first describe the observations and data reduction procedures. Then, we comment on the UV photometric properties of the galaxies (e.g., IRX-$\\beta$). We give detailed descriptions of each system, then discuss their detectability at high-$z$. Finally, we consider our systems in relation to Lyman-break galaxies, extremely red objects, and sub-mm galaxies. Here, we assume a Hubble Constant of $H_0=70$ km/s/Mpc, with $\\Omega_M=0.3$ and $\\Omega_\\Lambda=0.7$. Computing angular size and luminosity distances in a cosmology with nonzero $\\Lambda$ is described by \\citet{car92}. ", + "conclusions": "We have presented new UV images of a sample of seven LIGs and ULIGs taken with the HST. Our principal results are as follows. \\begin{enumerate} \\item{All seven of our ULIGs were detected in the UV. They show star clusters and extended emission, with the brightest UV light projected within the central kpc of the most IR-luminous members of the systems in 5/7 cases.} \\item{However, the UV peaks are displaced by at least few hundred pc from the peaks in {\\em I}-band and near-IR images, and so presumably the far-IR peaks as well. At most a few percent of the total UV light comes from the inner 500 pc, where the majority of the far-IR energy is generated.} \\item{Most nuclei are reddened. However, even after correction for dust reddening using the IRX-$\\beta$ correlation, the observed light even in large apertures is insufficient, by factors of 3--75, to account for the far-IR emission. When the UV light in 2 kpc diameter apertures is compared with the FIR emission, the deficits typically range from 2-4 orders of magnitude; hence, the reddening is insufficient to account for the large infrared-excesses. We conclude that the compact nuclear starbursts or AGN that dominate the bolometric energies of these galaxies are highly obscured in the UV.} \\item{All of the galaxies in this study have been previously noted as showing signs of some AGN activity, usually via Seyfert 2 or LINER spectra. However, AGN are not clearly the source of the high infrared excesses in most cases. In only two cases, IRAS 08572+3915 NW and IRAS 19254--7245 S, is it likely that AGN could dominate the bolometric energy output of the system; the majority appear to be dominated by star formation. In only one of the two possible AGN, IRAS 19254--7245 S, we may detect direct UV emission from the AGN. Even in this case, there is no emission in the FUV, and the NUV source is extended, suggesting that a circumnuclear disk, jet, or even star formation may be what is seen in the NUV.} \\item{Artifically redshifting the UV fluxes shows that six of our galaxies would be detectable in HDF-type exposures out to redshifts of at least $z=1.54$, yet cosmological surface brightness dimming would render all but two undetectable by $z=3.12$. The galaxies would also have sub-mm fluxes below current detection limits, if placed in that redshift range. Current optical and sub-mm surveys would therefore probably not detect most of our ULIGs, some of the most luminous objects in the local universe, if they were placed at redshifts much greater than 1.5.} \\item{We have compared our galaxies with 100 {\\em U}-band dropouts at $z\\approx 3$ from the HDF. If placed at $z\\sim 3$, most of our galaxies would be 2--3 magnitudes too faint for robust selection as {\\em U}-band dropouts in catalogs with HDF-like depths. Only one of our galaxies, the UV-bright member of VV 114, would have made it into the HDF sample. This suggests that ULIGs will not be found in samples of bright Lyman-break galaxies, though some could perhaps be lurking amongst the fainter and redder LBGs.} \\item{From estimating the $R$ and {\\em K}-band fluxes at $z=1.96$ and $z=3.80$ for the IR-dominant nuclei of our five ULIGs, we find that 2/5 of them would have colors very close to those of Extremely Red Objects. While some ULIGs may appear as EROs at high-$z$, in general, their faint but significant UV emission makes them somewhat too blue to be classified as EROs.} \\item{While our sample is small, the predicted faintnesses and red colors of our ULIGs, if placed at high redshift, are reminiscent of the extremely faint/red counterparts of many sub-mm galaxies. However, many sub-mm galaxies have counterparts brighter and/or bluer than our ULIGs would be at $z\\approx2$ or $z\\approx 3.8$. One sub-mm galaxy has been shown to have an extended radio source; if this is generally true, then that would call into question the use of ULIGs as low-redshift counterparts.} \\end{enumerate}" + }, + "0112/astro-ph0112487_arXiv.txt": { + "abstract": "We present evidence that low-mass starless cores, the simplest units of star formation, are systematically differentiated in their chemical composition. Molecules including CO and CS almost vanish near the core centers, where the abundance decreases by at least one or two orders of magnitude with respect to the value in the outer core. At the same time, the N$_2$H$^+$ molecule has a constant abundance, and the fraction of NH$_3$ increases toward the core center. Our conclusions are based on a systematic study of 5 mostly-round starless cores (L1498, L1495, L1400K, L1517B, and L1544), which we have mapped in C$^{18}$O(1--0), CS(2--1), N$_2$H$^+$(1--0), NH$_3$(1,1) and (2,2) and the 1.2 mm continuum (complemented with C$^{17}$O(1--0) and C$^{34}$S(2--1) data for some systems). For each core we have built a spherically symmetric model in which the density is derived from the 1.2 mm continuum, the kinetic temperature from NH$_3$, and the abundance of each molecule is derived using a Monte Carlo radiative transfer code which simultaneously fits the shape of the central spectrum and the radial profile of integrated intensity. Regarding the cores for which we have C$^{17}$O(1--0) and C$^{34}$S(2--1) data, the model fits these observations automatically when the standard isotopomer ratio is assumed. As a result of this modeling, we also find that the gas kinetic temperature in each core is constant at approximately 10 K. In agreement with previous work, we find that if the dust temperature is also constant, then the density profiles are centrally flattened, and we can model them with a single analytic expression. We also find that for each core the turbulent linewidth seems constant in the inner 0.1 pc. The very strong abundance drop of CO and CS toward the center of each core is naturally explained by the depletion of these molecules onto dust grains at densities of 2-6 $\\times 10^4$ cm$^{-3}$. N$_2$H$^+$ seems unaffected by this process up to densities of several $10^5$ or even $10^6$ cm$^{-3}$, while the NH$_3$ abundance may be enhanced by its lack of depletion and reactions triggered by the disappearance of CO from the gas phase. With the help of the Monte Carlo modeling, we show that chemical differentiation automatically explains the discrepancy between the sizes of CS and NH$_3$ maps, a problem which has remained unexplained for more than a decade. Our models, in addition, show that a combination of radiative transfer effects can give rise to the previously observed discrepancy in the linewidth of these two tracers. Although this discrepancy has been traditionally interpreted as resulting from a systematic increase of the turbulent linewidth with radius, our models show that it can arise in conditions of constant gas turbulence. ", + "introduction": "Dense molecular cores are the basic units of star formation in nearby clouds like Taurus and Perseus, where stars like our Sun have been forming over the last few million years \\citep[e.g.,][]{mye99}. The study of the physical structure and kinematics of these cores is therefore crucial for our understanding of the star formation process, and molecular lines play a role in every step of this work. They probe density and temperature through their excitation, and turbulent and systematic motions through their linewidth and Doppler shifts. For this reason, chemical anomalies in the core gas can hinder our attempt to understand core properties, as the lack of a full theory of core chemical composition has made it a standard practice to assume a homogeneous abundance for all molecular species. The presence of chemical inhomogeneities in the star-forming material at scales of dark clouds has been known for some time, with TMC-1 and L134N being the most studied examples \\citep[e.g.,][]{lit79,pra97,swa89,dic00}. The large-scale abundance gradients in these clouds seem best explained with time-dependent, gas-phase chemistry, implying that different condensations have evolved with different time scales (see \\citealt{lan00} and \\citealt{van98} for recent reviews). At the smaller size of the dense cores, a series of recent observations has shown that in some cases, the abundance of molecules like CO and CS decreases toward the core center (L1498: \\citealt{kui96,wil98}, IC5146: \\citealt{kra99,ber01}, L977: \\citealt{alv99}, L1544: \\citealt{cas99}, L1689B: \\citealt{jes01}). These decreases in abundance have been interpreted as resulting from the depletion of molecules onto dust grains at the high densities and cold temperatures occurring in dense core interiors (e.g., \\citealt{ber97,cha97}). It is not clear, however, whether these drops in abundance are typical of all dense cores, or they are limited to a small number of objects. To answer this question, we have carried out a systematic study of a sample of 5 starless cores by observing them in a similar manner and analyzing their emission with the same radiative transfer modeling. ", + "conclusions": "We have mapped 5 starless cores (L1498, L1495, L1400K, L1517B, and L1544) in C$^{18}$O(1--0), CS(2--1), N$_2$H$^+$(1--0), NH$_3$(1,1) and (2,2), C$^{17}$O(1--0) (3 cores), C$^{34}$S(2--1) (1 core), and the 1.2 mm continuum (4 cores), and complemented these data with a published 1.2mm continuum map of L1544 \\citep{war99}. For each core, we have self-consistently determined the radial profile of density (from the 1.2mm continuum), temperature (from NH$_3$), and molecular abundance (using a radiative transfer Monte Carlo model). To do this, we have fit simultaneously the central spectrum for each line and the radial profile of integrated intensity. As a result of this work, we conclude the following: 1. Each core is well fit with a radial density profile of the form $n(r) = n_0/(1+(r/r_0)^\\alpha)$ where $n_0$ ranges from about $10^5$ to $10^6$ cm$^{-3}$, $r_0$ is of the order of 3000-10000 AU, and $\\alpha$ ranges from 2 to 4. This type of profile naturally fits the central flattening and the large-$r$ power-law behavior found by previous authors. 2. All cores show evidence for a strong central drop in the abundance of C$^{18}$O and CS. This drop can be fitted with a negative exponential dependence on density ($\\exp (-n(r)/n_d)$), with an e-folding $n_d$ parameter of 2-6 $\\times 10^4$ cm$^{-3}$. C$^{17}$O(1--0) and C$^{34}$S(2--1) observations of several cores confirm these results and show that our radiative transfer models can naturally fit this rare isotopomer emission without extra free parameters. Although limited by the resolution of our data ($\\sim 50''$), we estimate that the C$^{18}$O and CS central abundance drops are at least of one or two orders of magnitude in each of the cores. 3. In contrast to C$^{18}$O and CS, the N$_2$H$^+$ abundance seems to be constant inside each core. The combination of narrow N$_2$H$^+$(1--0) central spectra and the presence of narrow CS(2--1) self absorptions indicate that the turbulent linewidth in each core is constant with radius. 4. The NH$_3$(1,1) and (2,2) lines are well fit with constant temperature models of 9.5 K (8.75 K in L1544) in which the NH$_3$ abundance increases toward the core center. The fact that the temperature is constant at densities for which gas and dust may be thermally coupled supports the assumption of constant dust temperature used in our 1.2mm continuum analysis. 5. The combination of a central C$^{18}$O and CS abundance drop, a constant N$_2$H$^+$ abundance, and a central NH$_3$ abundance increase seems to be a systematic characteristic of starless cores. This suggests that star-forming material is very chemically inhomogeneous before starting to collapse, and that searches for infall in starless cores should take this fact into account. 6. The pattern of chemical differentiation, especially for C$^{18}$O, CS, and N$_2$H$^+$ is in qualitative agreement with depletion models. Quantitative agreement between the data and these models can be achieved for certain choices of their parameters, which suggest the presence of polar ices on grains and a history of core contraction more rapid than predicted by subcritical ambipolar-diffusion models. These models, however, seem to overestimate N$_2$H$^+$ depletion at high densities and to miss the observed central NH$_3$ abundance increase. 7. The presence of chemical differentiation in cores automatically explains the systematic difference between map sizes of different tracers, a problem which has remained unexplained for more than a decade. Detailed Monte Carlo radiative transfer models show that CS maps, because of the central abundance decrease and radiative transfer effects, are expected to appear at least twice as large as NH$_3$ and N$_2$H$^+$ maps. These latter molecules are therefore more faithful tracers of the dense core material. 8. Comparing the CS and NH$_3$ linewidths from our Monte Carlo radiative transfer calculations, we find a systematic NH$_3$/CS linewidth discrepancy similar to that found by other authors \\citep[e.g.,][]{zho89}. These linewidth discrepancies have been traditionally interpreted as resulting from a systematic increase in the turbulent linewidth with radius, but our models show that it can arise from a combination of optical depth and self absorption in the CS line. This suggests that the linewidth-size relation does not apply in the inner 0.1 pc of cores." + }, + "0112/astro-ph0112164_arXiv.txt": { + "abstract": "We report the results of the spectral and timing analysis of observations of the Vela pulsar with the {\\sl Chandra} X-ray Observatory. The spectrum shows no statistically significant spectral lines in the observed 0.25--8.0 keV band. It consists of two distinct continuum components. The softer component can be modeled as either a magnetic hydrogen atmosphere spectrum with $kT_{{\\rm eff},\\infty} = 59\\pm 3$ eV, $R_\\infty =15.5\\pm 1.5$ km, or a standard blackbody with $kT_\\infty =129\\pm 4$ eV, $R_\\infty = 2.1\\pm 0.2$ km (the radii are for a distance of 250 pc). The harder component, modeled as a power-law spectrum, gives photon indices depending on the model adopted for the soft component: $\\gamma=1.5\\pm 0.3$ for the magnetic atmosphere soft component, and $\\gamma=2.7\\pm 0.4$ for the blackbody soft component. Timing analysis shows three peaks in the pulse profile, separated by about 0.3 in phase. Energy-resolved timing provides evidence for pulse profile variation with energy. The higher energy ($E > 1.8$ keV) profile shows significantly higher pulsed fraction. ", + "introduction": "The Vela pulsar (B0833-45), with a period of 89 millisecond and a characteristic age of 11,000 years, it is the prototype of the young ``Vela-like'' pulsars. The excellent angular resolution of {\\em Chandra} allows us to separate the pulsar from its synchrotron nebula and study its X-ray properties. \\vspace{-0.6cm} ", + "conclusions": "\\begin{itemize} \\item The HRC-S/LETG spectrum shows no conclusive evidence for lines in the 0.25--2.0 keV range. The spectrum shows two components (thermal + non-thermal), similar to middle-aged pulsars. The thermal component is consistent with the emission from a magnetic Hydrogen atmosphere. Extrapolation of the PL component matches the optical and the hard X-ray points. \\item The thermal luminosity of the Vela pulsar is lower than that predicted by the ``standard'' neutron star cooling models (Tsuruta 1998). \\item HRC data show three peaks with total pulsed fraction of about $9\\%$. Energy-resolved ACIS pulse profiles show energy-dependent shape and pulsed fraction. The pulse profile at $E < 1.8$ keV is similar to that observed with HRC. At $E > 1.8$ keV, the pulse profile shows only two peaks with estimated intrinsic pulsed fraction of $62\\pm20\\%$. \\end{itemize}" + }, + "0112/astro-ph0112022_arXiv.txt": { + "abstract": "It is shown that many properties of \\lya\\ forest absorbers can be derived using simple physical arguments. Analytical expressions are derived for the density and the size of an absorber as a function of its neutral hydrogen column density, which agree well with both observations and hydrodynamical simulations. An expression is presented to compute $\\Omega_{\\rm IGM}$ from the observed column density distribution, independent of the overall shape of the absorbers. Application to the observed column density distribution shows that at high redshift most of the baryons are in the forest and suggests a simple interpretation for its shape and evolution. ", + "introduction": "In the last decade semi-analytic models (e.g., \\cite{bi92:lya}) and hydrodynamical simulations (see \\cite{efstathiou00:lya} for a recent review) have been used to show that cold dark matter models with a nearly scale-invariant spectrum of initial, adiabatic, Gaussian fluctuations are very successful in reproducing the large body of high-quality observations of the \\lya\\ forest. The physical picture that has emerged is that the forest arises in a network of sheets, filaments, and halos, which give rise to absorption lines of progressively higher column densities. The low column density absorption lines ($N_{HI} \\lsim 10^{14.5}~\\cm^{-2}$ at $z\\sim 3$), arise in a smoothly fluctuating intergalactic medium (IGM) of moderate overdensity ($\\rho < 10\\left < \\rho \\right >$), and contain most of the baryons in the universe. On large scales the gas traces the dark matter and observations of the \\lya\\ forest can be used to study the large-scale distribution of matter, while on small scales hydrodynamical effects are important and the detailed line profiles can be used to reconstruct the thermal history of the IGM. In this contribution I will show that the modern picture of the forest can be derived directly from the observations using straightforward physical arguments, without making any assumptions about the presence of dark matter, the mechanism for structure formation or the precise cosmological model. This work is described in more detail in a recent publication \\cite{schaye01:lya}. ", + "conclusions": "" + }, + "0112/astro-ph0112036_arXiv.txt": { + "abstract": "Doppler spectroscopy of \\rhc\\ has detected evidence of a companion with an orbital period of 14.65 days and a minimum mass of 0.88 Jupiter masses. Astrometric observations performed with the Hubble Space Telescope Fine Guidance Sensor 1r using a novel new observing technique have placed an upper limit on the astrometric reflex motion of \\rhc\\ in a time period of only one month. These observations detected no reflex motion induced by the 14.65 day period radial velocity companion, allowing us to place a $3\\sigma$ upper limit of $\\sim$0.3\\,milliarcseconds on the semi-major axis of this motion, ruling out the preliminary Hipparcos value of 1.15\\,milliarcseconds. The corresponding upper limit on the true mass of the companion is $\\sim$30\\,\\mj, confirming that it is a sub-stellar object. ", + "introduction": "The Doppler spectroscopy technique has been used in recent years to detect low amplitude, periodic radial velocity variations in $\\sim$60 nearby stars, which have been interpreted as due to planetary mass companions. This interpretation requires several assumptions, namely that the root cause of the variation is Keplerian in nature; that the companion mass (\\mc) is substantially less than the mass of the primary; and that one is seeing light from a single star as opposed to an unresolved, comparable-mass binary system \\citep{impr98}. The quantity determined by the radial velocity observations is then \\mc\\,sin($i$) where $i$ is the unknown inclination of the orbital plane to our line of sight to the star. The argument usually advanced to suggest that the masses of the companions must be small assumes that the distribution of orbital inclinations is uniform so that, on average, $\\langle$\\mc$\\rangle\\,=\\,4\\langle$\\mc\\,sin($i)\\rangle/\\pi$ \\citep{chandra50}, which implies that the true companion mass cannot be much greater than the minimum mass, \\mc\\,sin($i$). Note that this statement is only true for a large sample, and it does not preclude individual companions from having true masses that are significantly larger than their minimum mass. Several pieces of work support the planetary mass interpretation. For example, \\citet{itomay01} have determined that the dynamical stability of the Ups And system constrains the mass of the outer planet to be less than 1.43 times its minimum \\mc\\,sin($i$) value. There is also one case, HD\\,209458, for which the companion has been observed to transit the disk of the central star \\citep{charb00}, hence giving a measure of both the companion radius and the inclination of its orbit. The companion mass of $\\sim$0.6 Jupiter masses (\\mj) is consistent with that of a giant planet. However, despite the statistical and dynamical arguments, and the special case of HD\\,209458, doubts have persisted that \\mc\\,sin(i) is nearly equal to \\mc\\ for several reasons. Using a combination of Hipparcos and ground-based MAP (Multichannel Astrometric Photometer) astrometry in conjunction with the orbital parameters derived from the radial velocity (RV) data, \\citet{gwetal00} have found that the \\mc\\,sin($i$) = 1\\,\\mj\\ companion to $\\rho$ CrB is in a nearly face-on orbit with $i$ = 0.5\\degr, and has a true mass of 115\\,\\mj, making it an M dwarf rather than a planetary-mass object. \\citet{zm00} have used Hipparcos astrometry to show that the \\mc\\,sin($i$) = 6.4\\,\\mj\\ companion to HD\\,10697 is actually a 38\\,\\mj\\ brown dwarf in an orbit with an inclination of just 5\\degr. A similar analysis combining the Hipparcos and radial velocity data extended to the 30 systems with orbital periods in excess of ten days \\citep{hanetal01} suggests that at least four of the 30 stars they analyzed have stellar mass companions, that is, \\mc\\,$>\\,80$\\,\\mj, which in turn would require $i<1$\\degr. If the distribution of inclinations in the sample is uniform, the probability of a system having an inclination $i\\,<\\,i_o$ is given by $1\\,-\\,$cos($i_o$). For $i_o$\\,=\\,1\\degr, the probability would be 1.5$\\times\\,10^{-4}$, making it unlikely that even one such system would be observed in a sample of 2000 stars. This led Han et al. to suggest that there might be a bias toward small inclination angles in the radial velocity studies. \\citet{pour01} also finds statistically significant astrometric orbits with low inclinations for three out of four of the Han et al.~potential stellar mass companions. However, \\citet{pourar01} argue that the trend to low inclinations is an artifact of the adopted reduction procedure, and that the astrometric data are not precise enough to allow the conclusion that a significant fraction of the RV companions have stellar masses. An independent line of evidence also raises the possibility of stellar mass companions for some of the 60 systems. Among a subset of 9 of the stars of spectral type F with candidate planetary companions, Suchkov \\& Schultz (2001) have identified three potential binaries, HD\\,19994, HD\\,89744, and HD\\,169830, along with HD\\,114762 (private communication), based on the fact that they are overly bright for their spectral type and Hipparcos distance. Gonzalez et al. (2001) found that two stars with RV companions, HD\\,37124 and HD\\,46375, are similarly too luminous for their spectral type and distance; they might each be unresolved binary systems. However, they also note that a wide, long-period binary cannot be ruled out, so the short period RV companion need not be the source of the ``excess\" luminosity. Another line of reasoning suggesting that the RV companions may not be planetary in nature comes from an analysis of the distribution of their eccentricities and orbital periods, which are statistically indistinguishable from those for single-line spectroscopic binaries (SB1s) \\citep{step01,heacox99}. Moreover, the apparent correlation of eccentricity with orbital period for the RV companions is similar to that for SB1s \\citep{black97,heacox99}, and their bivariate probability distribution functions are again statistically indistinguishable (Stepinski \\& Black 2001). The observed orbital properties also differ strongly from those of our own planetary system. The current data, interpreted as planetary systems with random inclinations, also does not lead to a simple theoretical picture (e.g., Marcy et al.~1999), and there is currently no compelling dynamical argument to support the interpretation that all of the RV companions are planets, rather than brown dwarfs or stars in low inclination orbits. A determination of which companions (if any) are not planetary would facilitate the development of a dynamical model. Astrometric observations of the stellar reflex motion induced by a companion can potentially remove the uncertainty in sin($i$), and thus determine the companion mass. The Hubble Space Telescope (HST) Fine Guidance Sensor 1r (FGS1r) can measure relative stellar positions to an accuracy $\\sim$0.3\\,milliarcsecond (hereafter mas), a factor of 3-5 improvement over that of Hipparcos and MAP data. In this paper we present the results of a pilot program utilizing an observing technique designed to quickly search for perturbations larger than about 0.3\\,mas to the position of $\\rho^1$\\,Cnc (HR\\,3522, HD\\,75732, 55\\,Cnc), a G8V star with an \\mc\\,sin($i$) = 0.88\\,\\mj\\ companion in an orbit with a period of 14.65 days \\citep{betal97}. Han et al. report a preliminary reflex motion with a semi-major axis of 1.15\\,mas for this star. We discuss our rationale for selecting this target and our observing strategy in section 2, and our results in section 3. ", + "conclusions": "" + }, + "0112/astro-ph0112170_arXiv.txt": { + "abstract": "We have used the Hubble Space Telescope Wide--Field and Planetary Camera 2, in combination with ground--based spectroscopy, to measure the integrated flux of galaxies at optical wavelengths --- the extragalactic background light (EBL). We have also computed the integrated light from individual galaxy counts in the images used to measure the EBL and in the Hubble Deep Field. We find that flux in galaxies as measured by standard galaxy photometry methods has generally been underestimated by about 50\\%, resulting from missed flux in the outer, lower surface--brightness parts of galaxies and from associated errors in the estimated sky level. Comparing the corrected, integrated flux from individual galaxies with our total EBL measurement, we find that there is yet further light that contributes to the background that is not represented by galaxy counts, and that the total flux in individually detected sources is a factor of 2 to 3 less than the EBL from 8000 to 3000 \\AA. We show that a significant fraction of the EBL may come from normal galaxies at $z<4$, which are simply undetectable as a result of K-corrections and cosmological surface brightness dimming. This result is consistent with results from recent redshift surveys at $z<4$. In the context of some simple models, we discuss the constraints placed by the EBL on evolution in the luminosity density at $z>1$; while significant flux comes from galaxies beyond the current detection limits, this evolution cannot be tightly constrained by our data. Based on our measurements of the optical EBL, combined with previously published measurements of the UV and IR EBL, we estimate that the total EBL from 0.1--1000 \\micron\\ is 100$\\pm$20 nW m$^{-2}$ sr$^{-1}$. If the total EBL were produced entirely by stellar nucleosynthesis, then we estimate that the total baryonic mass processed through stars is $\\Omega_\\ast = 0.0062 (\\pm 0.0022) h^{-2}$ in units of the critical density. For currently favored values of the baryon density, $\\Omega_{\\rm B}$, this corresponds to $0.33\\pm0.12 \\Omega_{\\rm B}$. This estimate is smaller by roughly 7\\% if we allow for a contribution of 7$h_{0.7}$ nW m$^{-2}$ sr$^{-1}$ to the total EBL from accretion onto central black holes. This estimate of $\\Omega_*$ suggests that the universe has been enriched to a total metal mass of $0.21(\\pm0.13) Z_\\odot \\Omega_{\\rm B}$, which is consistent with other observational estimates of the cumulative metal mass fraction of stars, stellar remnants, and the intracluster medium of galaxy clusters in the local universe. ", + "introduction": "\\label{intro} The integrated optical flux from all extragalactic sources is a record of the stellar nucleosynthesis in the universe and the chemical evolution which has resulted from it. In \\BFM\\ 2001a (henceforth, Paper I), we presented detections of the optical EBL in the HST/WFPC2 wide-band filters F300W ($U_{300}$, $\\lambda_0\\sim$ 3000\\AA), F555W ($V_{555}$, $\\lambda_0\\sim$ 5500\\AA), and F814W ($I_{814}$, $\\lambda_0\\sim$ 8000\\AA) based on simultaneous data sets from Hubble Space Telescope (HST) and Las Campanas Observatory (LCO). In \\BFM\\ 2001b (henceforth, Paper II), we presented details of a measurement of the diffuse foreground zodiacal light which we use in Paper I. Here we briefly summarize the results of Papers I and II and discuss the cosmological implications of these detections of the EBL. The majority of the EBL at UV to IR wavelengths is produced by stars at restframe wavelengths of 0.1--10\\micron. Due to cosmic expansion, the EBL at $U_{300}$, $V_{555}$, and $I_{814}$ potentially includes redshifted light from stellar populations out to $z\\sim8$ (the redshifted Lyman--limit cut--off of the $I_{814}$ filter). Although stars themselves do not emit much light at wavelengths longer than 10\\micron, a complete census of the energy produced by stellar nucleosynthesis in the universe must consider the EBL over the full wavelength range 0.1-1000\\micron\\ because dust in the emitting galaxies will absorb and re-radiate starlight, redistributing energy from nucleosynthesis into the thermal IR. With 8m--class telescopes and HST, the limits of resolved--source methods (i.e., number counts, redshift surveys, QSO absorption lines, etc.) for studying star formation in the universe are being extended to ever fainter levels; however, a direct measurement of the EBL remains an invaluable complement to these methods. Galaxies with low {\\it apparent} surface brightness --- both intrinsically low surface brightness galaxies at low redshift and normal surface brightness galaxies at high redshift --- are easily missed in surface--brightness--limited galaxy counts and consequentally in follow--up redshift surveys. Identification, not to mention photometry, of faint galaxies becomes very uncertain near the detection limits. Even efforts to understand galaxy evolution, chemical enrichment, and star formation through QSO absorption line studies appear to be biased against chemically enriched, dustier systems, as these systems can obscure QSOs which might lie behind them (Fall \\& Pei 1989, Pei \\& Fall 1995, Pettini \\etal 1999). In contrast, a direct measurement of the spectral energy distribution (SED) of the EBL from the UV to the far--IR is a complete record of the energy produced by star formation and is immune to surface brightness selection effects. In addition to energy originating from stellar nucleosynthesis, the EBL includes energy emitted by accreting black holes in quasars and active galactic nuclei. However, at optical wavelengths, the quasar luminosity functions at redshifts $z\\lesssim5$ indicate that the optical luminosity density from quasars is a small fraction ($\\sim2.5$\\%) of the that from galaxies (e.g. Boyle \\& Terlevich 1998). In addition, our measurement of the EBL excludes any point--like sources (of which there are 3 in our images), under the prior assumption that those sources are Galactic foreground stars. We therefore expect quasars to be a negligible source of flux in our measurements of the optical EBL. The contribution from active galactic nuclei (AGN) is more difficult to assess, as recent dynamical evidence (Richstone \\etal 1998) indicates that massive black holes reside in most galaxies and sensitive optical spectroscopy (Ho \\etal 1997a, 1997b) indicates that AGN have at least a weak contribution to more than 50\\% of nearby galaxies. Nonetheless, simple accretion models, the total X-ray background, and the X-ray to far--IR spectral energy distribution of AGN and quasars all indicate that the total contribution to the bolometric EBL from accretion onto central black holes is $\\lta15$\\% (see \\S\\ref{agn}), and is emitted at thermal IR wavelengths. In principle, measurements of the EBL also constrain possible the total energy output from more exotic sources, such as gravitationally collapsing systems, brown dwarfs, and decaying particles (see Carr, Bond, \\& Hogan 1986, 1991 and Dwek \\etal 1998 for discussions). The outline of the paper is as follows. In \\S\\ref{sum.obs}, we give an overview of the observations and methods used to measure the EBL as discussed in Papers I and II. In \\S\\ref{sum.detect}, we summarize the individual measurements and associated errors we have obtained from each data set and the final EBL detections which result from them. In \\S\\ref{opt.res}, we compare the measured EBL with the integrated optical flux from resolvable sources as quantified by number counts and luminosity functions. In \\S\\ref{opt.unres}, we quantify the contributions to the optical EBL which one might expected from sources which fall below the detection limits of the HDF based on explicit assumptions regarding the surface brightness, luminosity, and redshift distribution of galaxy populations in the universe. In \\S\\ref{bolom}, we discuss models of the SED of the EBL based on these and recent results in the far infrared. Finally, in \\S\\ref{stel} we discuss the total star formation and chemical enrichment history of the universe required to produce the bolometric flux of the EBL, and compare the inferred values to other observations of the total baryon fraction in stars and the metal mass density in the local universe. We abbreviate the adopted units \\escsa\\ as cgs throughout. ", + "conclusions": " \\noindent (1) We find that the corrected number counts at $V$ and $I$ magnitudes fainter than 23 \\ABmag\\ obey the relation $N\\propto 10^{\\alpha m}$ with $\\alpha= 0.33\\pm0.01$, and $\\alpha= 0.34\\pm0.01$, respectively, which is consistent with the slope found at brighter magnitudes (e.g. Smail \\etal 1995, Tyson 1988). This is significantly steeper than the slope of the raw HDF number counts ($\\sim 0.24\\pm 0.1$ at $V>23$ \\ABmag, and $\\sim 0.22\\pm 0.1$ at $I>23$ \\ABmag). In contrast with the raw counts, the corrected counts show no decrease in slope to the detection limit. If we integrate the corrected number counts down to an apparent magnitude corresponding roughly to a dwarf galaxy ($M_V\\sim -10$ mag) at $z\\sim3$, $V\\sim38$ \\ABmag, we obtain a total flux of $1.2$\\tto{-9} cgs in both $V$ and $I$. This is $1.2\\sigma$ below the mean EBL23 flux we estimate at $V_{606}$ ($I_{814}$), suggesting that number counts would need to be steeper over some range in apparent magnitude fainter than the current detection limits in order to obtain the mean EBL flux we detect, or that the value of EBL23 is roughly $\\sim1\\sigma$ below our mean detections at $V$ and $I$. \\noindent (2) Based on a local luminosity density consistent with Loveday \\etal (1992), passive evolution in the luminosity density of galaxies under-predicts the EBL by factors of roughly 3, 2, and 2 at $U_{300}$, $V_{555}$, and $I_{814}$, respectively. Note, however, that if the local luminosity density is a factor of two higher than the Loveday \\etal values we have adopted here, as found by Blanton \\etal (2001), then passive evolution agrees with the flux in resolved galaxies (minEBL23) and with our mean EBL detections to within $1\\sigma$. The mean detected EBL therefore requires stronger evolution in the luminosity density than passive evolution will produce, however, the exact form of that evolution is not well constrained by our results. Adopting the local luminosity density assumed by Lilly \\etal (1996, CFRS), the $1\\sigma$ upper limits of the cumulative flux measured by Lilly \\etal from redshifts $01$. If we adopt ${\\cal L}(\\lambda,z)\\propto (1+z)^{\\delta(\\lambda)}$ for the luminosity density at $01$, such as suggested by Steidel \\etal (1999) is consistent with the detected flux in sources at $V_{555}$ and $I_{814}$, and with the detected EBL at $U_{300}$. At the upper limit of the EBL detections, we find that the luminosity density can continue to rise as a power law to $z\\sim 2.5$ without over--predicting the EBL. \\noindent (3) We have modeled the effects of cosmological K-corrections, passive evolution, and $(1+z)^4$ cosmological surface brightness dimming on the detectability of local--type galaxy populations as a function of redshift. For these models, we have adopted the spatial resolution and surface brightness limits of the HDF. For models which bracket the observed surface brightness distribution of galaxies in the local universe, we find that roughly 10--40\\% of the EBL from galaxies fainter than $V\\sim23$ (i.e. those sampled in an HDF--sized image), comes from galaxies which are, at present, individually undetectable at wavelengths $\\lambda>4500$\\AA, and roughly 20--70\\% comes from individually undetected galaxies at $\\lambda<4500$\\AA. Most of the flux from a local-type galaxy population located at $z=3$ would come from sources that would not be individually detected in the HDF. Our models indicate that the true EBL is likely to be between the mean detected EBL23 values and the $1\\sigma$ lower limits of those detections at $V$ and $I$, and within $\\pm 1\\sigma$ at $U$. \\noindent (4) Scaling the model of the bolometric EBL derived by Dwek \\etal (1998), which is based on a combined UV-optical estimate of the star formation rate and a model for dust obscuration and re--emission based on the spectrum of IRAS sources, we find that the optical EBL we detect corresponds to a total bolometric EBL (0.1 to 1000\\micron) of 100$\\pm20$ nW~m$^{-2}$sr$^{-1}$. \\noindent (5) From this estimate of the total bolometric EBL, we estimate that the total baryonic mass processed through stars is $\\Omega_\\ast = 0.0062(\\pm0.0012) h^{-2}= 0.33(\\pm0.07) \\Omega_{\\rm B}$, and that the mean metal mass density in the universe is $\\Omega_Z= 0.0040(\\pm0.0022) h^{-2} Z_\\odot = 0.24(\\pm0.13) Z_\\odot \\Omega_{\\rm B}$, for $\\Omega_{\\rm B}=0.019(\\pm 0.001)h^{-2}$ (Burles \\& Tytler 1998). These estimates are consistent with limits from other observational constraints." + }, + "0112/astro-ph0112346_arXiv.txt": { + "abstract": "Multi-fiber spectroscopy has been obtained for 335 galaxies in the field of the double cluster A3128/A3125, using the 2dF multi-fiber positioner on the Anglo-Australian Telescope. When combined with previously published results, a total of 532 objects in the double cluster now have known redshifts. We have also obtained a 20 ks {\\it Chandra} ACIS-I image of the central 16\\arcm x 16\\arcm of A3128 and radio imaging of the cluster with the Molonglo Observatory Synthesis Telescope and the Australia Telescope Compact Array. The spatial-kinematic distribution of redshifts in the field of A3128/A3125, when combined with the {\\it Chandra} ACIS-I image of A3128, reveals a variety of substructures present in the galaxy distribution and in the hot intracluster medium (ICM). The most striking large-scale feature in the galaxy distribution is a relatively underpopulated redshift zone $\\sim$4000 \\kms on either side of the mean cluster velocity at $\\sim$17500 \\kms. We attribute this depletion zone to the effect of the extensive Horologium-Reticulum (H-R) Supercluster, within which A3128/A3125 is embedded. In addition to this large-scale feature, numerous smaller groups of galaxies can be identified, particularly within the underpopulated region within $\\pm$4000 \\kms of the mean cluster redshift. Due to the large gravitational influence of the H-R Supercluster, these groups arrive at A3128 with a high infall velocity, well in excess of the local sound speed. Two of these groups appear as elongated filaments in position-velocity diagrams, indicating that they are tidally distended groups which have been disrupted after a close passage through A3128. In fact, A3125 itself appears to be in such a post passage condition. We have identified a primary NE-SW merger axis connecting A3128 with A3125, along which the filaments are also oriented. In addition, the {\\it Chandra} image reveals that the X-ray emission is split into two components, each with very small core radii, that are separated by $\\sim$1 Mpc along the NE-SW axis. We have combined the redshift, X-ray, and radio data to propose that the complex X-ray morphology revealed in the {\\it Chandra} image is likely the result of a hypersonic infall of a relatively small group into A3128. The group produces a major disruption in the ICM due to its high infall velocity. ", + "introduction": "The distribution of groups, clusters, and superclusters of galaxies represents a fundamental testing ground for theories of the origin and evolution of structure in the universe. Until recently, the intersection between observation and theory has largely relied on statistical properties of galaxy spatial/kinematic clustering, along with X-ray determined global ICM temperatures and azimuthally averaged radial brightness and temperature profiles. In contrast, much of the existing observational data indicates that large-scale asymmetric, probably filamentary, structures, which are not easily subjected to statistical definition {\\it on an individual basis}, are present in clusters and superclusters (e.g., Gregory \\& Thompson 1978; Shandarin 1983; de Lapparent, Geller, \\& Huchra 1986; West, Jones, \\& Forman 1995; West \\& Blakeslee 2000). Multi-fiber spectroscopy of galaxies with 400-fiber positioners deployed over 2\\arcdeg~ fields, as well as the unprecedented combination of spatial resolution and sensitivity in X-rays provided by the {\\it Chandra} and XMM-Newton Observatories, are rapidly advancing the observational view of clusters of galaxies. In addition, numerical simulations, carried out within the framework of a cold dark matter dominated universe in which structure is built up in a hierarchical fashion, have reecently reached the point where large areas, on scales comparable to the largest superclusters, can be simulated in some detail (Pearce \\etal 2001 and references therein). Furthermore, using parametrized physics of the baryonic component, simulations within the framework of a cold dark matter dominated universe now predict the time evolution of the baryonic material (Kauffmann \\etal 1999a,b; Somerville \\& Primack 1999; Benson \\etal 2001 and references therein), thus providing a more direct link to the observational data. Hence the opportunity now presents itself to make a comprehensive comparison between observations and simulations in the case of {\\it specific} clusters/superclusters, thereby providing a more sensitive probe of structure evolution than can be gained from purely statistical considerations. Recent studies of the Shapley Supercluster, an extremely massive concentrations of galaxies and clusters first pointed out by Shapley (Shapley 1930; Raychaudhury 1989), represent one example of the possibilities provided by an intensive campaign of multi-fiber spectroscopy, as well as of X-ray and radio imaging. These studies have produced a wealth of data on the spatial/kinematic structures present in the Shapley Supercluster among the galaxies (e.g., Quintana \\etal 2000; Drinkwater \\etal 1999; Bardelli, Zucca, \\& Baldi 2001) and on the hot ICM (e.g., Ettori \\etal 2001), as well as results on the the emission line and radio properties of cluster galaxies (Baldi, Bardelli, \\& Zucca 2001; Venturi \\etal 2001). In the process, a variety of structures have been found, and the influence of environment on both galaxy and radio source evolution has been partially elucidated. In this paper we present new observational data on a double galaxy cluster, A3128/A3125, which is itself embedded in the massive Horologium-Reticulum (H-R) Supercluster. As is discussed later, the H-R Supercluster, and the above-mentioned Shapley Supercluster, represent the two largest cluster concentrations within the local 300 h$_{50}^{-1}$ Mpc\\footnote{Throughout this paper we adopt $H_0 = 50\\,h_{50}^{-1}$ \\kms}. The observational data consist of multi-fiber spectroscopy of the double cluster A3128/A3125 with the Anglo-Australian Telescope and 2\\arcdeg~field multi-fiber positioner. In addition, we present X-ray imaging of the the central regions of A3128 and A3125 with the Chandra X-ray Observatory. Radio observations obtained with the Molonglo Observatory Synthesis Telescope and the Australia Telescope Compact Array are also presented. These observations have revealed a variety of substructures, both in galaxy groupings and in the multiple X-ray peaks detected, which in responding to the large gravitational acceleration provided by the H-R Supercluster, reach unusually high infall velocities into A3128. Thus the connection between A3128/A3125 and the surrounding H-R Supercluster provides an opportunity to characterize substructure merging on a variety of scales and unusual conditions. In \\S 2 we present the multi-fiber spectroscopy and X-ray and radio data. In \\S 3 we discuss the kinematic structures revealed by the multi-fiber spectroscopy, while in \\S 4 we discuss the observed state of the ICM, as inferred from the Chandra observations, and in \\S 5 we present the MOST and ATCA radio imaging. A discussion of our results is given in \\S 6; specifically, we attempt to produce a unified view of the merging events occurring in A3128/A3125 based on the full set of observations. ", + "conclusions": "We have collected optical redshift data for 532 objects in the field of the merging double cluster galaxies A3128/A3125. The redshift information has been supplemented by both X-ray and radio imaging. The goal of this program is to characterize and understand the variety of structures still present in the double cluster, and its relation to the larger environment of the Horologium-Reticulum Supercluster. Our principal conclusion is that a large number of substructures are still present in the A3128/A3125 system. This conclusion is based both on groups and filaments evident in the position-velocity space of the galaxy distribution and on the multi-component nature of the X-ray emitting ICM. The most striking large-scale feature in the galaxy distribution is the relatively low number density of galaxies for $\\sim$4000 \\kms on either side of the mean cluster velocity at $\\sim$17500 \\kms. We interpret this feature as due to the large gravitational potential well of the H-R Supercluster. Within the 4000 \\kms ``depleted'' zone, those galaxies that are present appear to be members of small groups or filaments. The latter are extended features in position-velocity space whose statistical reality is somewhat uncertain. We ascribe the filaments to the tidal distension of a group after it has fallen through the A3128/A3125 system and is emerging out the other side. In fact, A3125 itself shows some characteristics of a tidally disturbed system. We note that the filamentary structures appear to follow the main NE-SW axis along which most merging activity appears to be taking place. This is the axis which connects A3125 to the larger A3128 (at a projected separation of $\\sim$6 Mpc, and also along which the X-ray emission is split into two bright components, at a projected separation of $\\sim$1 Mpc. We have proposed two ongoing merging events to explain the current state of the A3128/A3125 system. The first is the merger event between A3125 and A3128, which is most plausibly explained if A3125 has already passed through A3128, and is highly dispersed as a result of the passage. The second involves the double peaked nature of the X-ray emission. Here we have noticed the close spatial correspondence between one particular high velocity filament and group with the X-ray emitting gas. Although the group/filament system is not massive, in comparison with the main body of A3128, the high infall velocity generated by the potential well of the H-R Supercluster produces a large energy deposition in the collision between the filament and A3128, and allows for only a modest infalling system to produce a major impact on the cluster ICM. Specifically, we argue that the morphology of the NE X-ray peak, along with its coincidence with the higher velocity end of the galaxy filament, indicates that the NE X-ray component represents the surviving ICM of the galaxy filament that has endured a hypersonic ($\\sim$Mach 6) encounter with A3128. The SW X-ray component appears to be the still intact ICM of a compact group that represents the still-infalling end of the galaxy filament-group. This gas is believed to be responding to the potential well of a compact group of galaxies surrounding a bright galaxy. While the details of this picture are still quite uncertain, the key ingredient is the high encounter velocity produced by the H-R Supercluster, which makes even the infall of a small group an energetic and interesting event in the life of a galaxy cluster. Further observations which could clarify this picture include deeper X-ray observations, to characterize the temperature structure in the ICM, deep high resolution radio imaging to clarify the nature of the putative radio arc associated with the NE X-ray peak, redshifts of galaxies in the apparent compact group that is coincident with SW X-ray component, and further redshift surveys throughout the H-R Supercluster." + }, + "0112/astro-ph0112420_arXiv.txt": { + "abstract": "The propagation of a shock wave into an interstellar medium is investigated by two-dimensional numerical hydrodynamic calculation with cooling, heating and thermal conduction. We present results of the high-resolution two-dimensional calculations to follow the fragmentation due to the thermal instability in a shock-compressed layer. We find that geometrically thin cooling layer behind the shock front fragments into small cloudlets. The cloudlets have supersonic velocity dispersion in the warm neutral medium in which the fragments are embedded as cold condensations. The fragments tend to coalesce and become larger clouds. ", + "introduction": "The interstellar medium (ISM) and cold gas clouds are characterized by a clumpy substructure and a turbulent velocity field (Larson 1981). The maintenance and dissipation processes of the turbulence are supposed to be important in the theory of star formation (McKee 1989, Nakano 1998, Elmegreen 1999). The understanding of the origin of cold clouds and their internal substructure has therefore fundamental importance for a consistent theory of star formation and ISM. Some of the polarization maps and direct measurements of field strength in some star forming regions suggested the importance of MHD (Alfv\\'{e}n waves) turbulence (Myers \\& Goodman 1988, Crutcher et al. 1993). Recent simulations of MHD turbulence, however, suggest that it dissipates rapidly (Gammie \\& Ostriker 1996, Mac Low et al. 1998, Mac Low 1999, Ostriker et al. 1999). Possible sources of energy supply are winds and outflows from young stellar objects (Franco \\& Cox 1983, McKee 1989). Note however that clumpy structures with supersonic velocity dispersions are also observed even in regions where star formation is inactive. Thus, the origin of clumpy cloud structure remains as an outstanding issue. We propose that the clumpiness in clouds arises naturally from their formation through thermal instability which acts on timescales that can be much shorter than the duration of the interstellar shocks (e.g., galactic spiral shocks and supernova remnants). The basic properties of the thermal instability are studied in a pioneering paper by Field (1965). Schwarz, McCray, and Stein (1972) studied numerically the growth of condensation in cooling region including the effects of ionization and recombination. Hennebelle \\& Perault (1999) studied the elementary condensation process in turbulent flow in the restricted conditions of neutral atomic gas in plane-parallel geometry. Burkert \\& Lin (2000) studied cooling and fragmentation of gas using simplified power-law cooling function. Smith (1980) studied collisions between cold atomic clouds, which produce thick layers of shock-heated atomic gas in which thin sheets of cold molecular gas form by the thermal instability. Koyama \\& Inutsuka (2000, hereafter Paper I) have done one-dimensional hydrodynamic calculations for the propagation of a strong shock wave into warm neutral medium (WNM) and cold neutral medium (CNM) including detailed thermal and chemical processes. They have shown that the post-shock region collapses into a cold layer as a result of the thermal instability. They expect that this layer will break up into very small cloudlets which have different translational velocities. In this paper, we show that the fragmentation of the shock-compressed layer indeed provides turbulent condensations, by using two-dimensional hydrodynamic calculation with radiative cooling and heating and thermal conduction. ", + "conclusions": "\\subsection{Formation of Molecules} The main coolant in cold dense clouds is supposed to be CO (see, e.g., Neufeld, Lepp, \\& Melnick 1995; Paper I). H$_2$ enables the formation of CO and other molecules. The abundance of H$_2$ in the CNM critically depends on the optical depth and local density. The typical column density of the layer becomes $2\\times 10^{20} {\\rm cm^{-2}}$. We calculated detailed thermal-chemical equilibrium state in the dense region by adopting the values of density, temperature, and column density from this dynamical simulation. Twenty-five percent of the hydrogen is H$_2$ and 0.03\\% of the carbon is CO in thermal and chemical equilibrium. We simulate the observation of molecular clumps at $t=1.06$ Myr (Figure \\ref{fig1}b) as the Position-Velocity (P-V) diagram of the ${}^{12}$CO J = 1 -- 0 emission (Figure \\ref{fig2}a). We use the usual relationship $T_{\\rm B}=\\frac{1}{2}\\lambda^2I/k_{\\rm B}$, with the specific intensity $I=n^2\\Lambda L/(4\\pi \\Delta v)$, where $n^2\\Lambda$ is the cooling per unit volume, $L$ is the path length, $\\Delta v$ is the Doppler line width. We adopt $\\Delta v$ to be the sound velocity. As shown in Figure \\ref{fig2}a, the typical velocity dispersion in the clumps is about a few km/s. The cloud to cloud velocity dispersion is also a few km/s. Figure \\ref{fig2}b shows that the P-V diagram of hydrogen nuclei. The velocity dispersion of the diffuse warm medium is about 10 km/s. \\subsection{An Origin of ``Turbulence'' in Interstellar Clouds} Diffuse ISM in the Galaxy is frequently compressed by supernova explosions (McKee \\& Ostriker 1977). Therefore, the shock propagation into the ISM plays an important role in the evolution of the Galactic ISM. We have studied the shock propagation into WNM by two-dimensional hydrodynamic calculations, and have shown that the thermally collapsed layer breaks up into very small cloudlets. Fragmentation of the thermally collapsing layer is a result of the thermal instability. The thermal instability produces many cloudlets which have different translational velocities. We expect that the Galactic ISM is occupied by these small cloudlets which have supersonic velocity dispersion , because the ISM is frequently compressed by supernova explosions, stellar winds, spiral density waves, cloud-cloud collisions, etc. The \\ion{H}{1} 21-cm observation maps show the existence of many shell-like or filamentary structures in the Galaxy (Hartmann \\& Burton 1997). In addition, high resolution observations suggested that the clumpy distribution is ubiquitous in the Galaxy (Heiles 1997). If these structures are really the results of the shock waves, many small molecular cloudlets should be hidden in the shock-compressed layers. These cloudlets should have translational velocities owing to the fragmentation of thermal instability. Observational ``turbulence'' in the ISM should reflect these motions. Thus, we propose that an origin of ``turbulence'' in the ISM is the motion of this small cloudlet complex. In the radiative shocked layers, the gases lose thermal energy through radiative cooling. Thus, the initial kinetic energy of the pre-shock gas (in the comoving frame of the post-shock gas) converted to the radiation energy, which escapes from the system. If, however, the post-shock gas becomes dynamically unstable by the thermal instability as in this paper, some portion of the thermal energy is transformed into the kinetic energy of the translational motions of the cold cloudlets, which does not easily escape from the system. Therefore we can consider the origin of interstellar turbulence is due to the conversion of the gas energy in supersonic motion. In principle, the kinetic energy of the cloudlets can be lost via coalescence of the cloudlets. The damping rate of the velocity dispersion of the cloudlets will play a key role in the evolution of this system, and hence, should be studied in the subsequent paper. Expected effects of magnetic field modify the development of the thermal instability described here. The presence of magnetic field can suppress thermal conduction efficiently, allowing the collapse of small scale structure by the thermal instability. The magnetic forces suppress the motion that is perpendicular to the magnetic field lines, and hence suppress the growth of perturbations whose wavevector is perpendicular to the magnetic field lines. However, the magnetic forces do not affect the motion that is parallel to the magnetic field lines, and hence the growth of perturbations whose wavevector is parallel to the magnetic field lines. As a result, the dense sheets or filaments will form and tend to align {\\em perpendicular} to the magnetic field lines. We need three-dimensional calculation to analyze these effects of magnetic fields. Numerical computations were carried out on VPP300/16R and VPP5000/48 at the Astronomical Data Analysis Center of the National Astronomical Observatory, Japan." + }, + "0112/astro-ph0112085_arXiv.txt": { + "abstract": "The structure of accretion discs around magnetic T Tauri stars is numerically calculated using a particle hydrodynamical code, in which magnetic interaction is included in the framework of King's dimagnetic blob accretion model. Setting up the calculation so as to simulate the density structure of a quasi-steady disc in the equatorial plane of a T Tauri star, we find that the central star's magnetic field typically produces a central hole in the disc and spreads out the surface density distribution. We argue that this result suggets a promising mechanism for explaining the unusual flatness (IR excess) of T Tauri accretion disc spectra. ", + "introduction": "It is now quite widely accepted that in the formation process of low-mass ($M \\la 2 M_\\odot$) stars, a disc accretion phase is typically present (see Hartmann 1998 for an extensive review and references). In particular, the classical T Tauri stars (CTTS) are thought to be still surrounded by discs and the emission properties of some typical objects suggest that these discs are not merely ``passive\" dust bodies, irradiated by the central star, but are rather bona-fide energy producing, viscous accretion discs, similar to the ones present in mass transferring close binary systems, and whose modeling dates back to the classical works of Shakura \\& Sunyaev (1973) and Lynden-Bell \\& Pringle (1974). Because of their ubiquity and of some interesting theoretical challenges they pose, accretion discs have remained at the centre stage of astrophysical research (see e.g. Frank, King \\& Raine 1992; Papaloizou \\& Lin 1995; Lin \\& Papaloizou 1996 for reviews). One particular issue in this context is the fact that in many CTTS the measured spectrum presents an unusually high emission in the infra red (called IR excess), far more than expected from disc reprocessing of stellar light or viscous heating of the disc (see Bertout 1989, Beckwith et al. 1990, Hartmann et al. 1994). Our wish to investigate this problem, as well as some other facets of CTTS discs, has been the main motivation for the work reported on in this paper. It consists of setting up a numerical calculation of the properties of an accretion disc around magnetic CTTS, using a scheme of magnetic interaction, originally proposed in the context of diamagnetic accretion by King \\shortcite{king1}. This description of the magnetic interaction assumes that as material moves through the magnetosphere it interacts with the local magnetic field via a velocity dependent acceleration (force per unit mass, $f_{\\rm mag}$) of the form: \\begin{equation} f_{\\rm mag} = - K \\, [{\\bf u} - {\\bf u}_f]_\\bot, \\label{general} \\end{equation} where ${\\bf u}$ and ${\\bf u}_f$ are the velocities of the material and magnetic field lines respectively, $K$ is a suitable \"magnetic drag\" coefficient (see below) and the suffix $\\perp$ refers to the vector component perpendicular to the field line. The magnetic acceleration, as expressed in equation \\ref{general}, is intended to represent the dominant term of the magnetic interaction, with $K$ containing the relevant parameters determining the effective magnetic time-scale. Within the above description (\\ref{general}) there are still a number of different possibilities to model the inner disc - magnetosphere interaction. In this paper we shall use the diamagnetic blob accretion (DBA) model, although it is possible to formulate the diametrically opposite case (complete magnetic penetration of the disc) in the general form (\\ref{general}) as well. Wynn, Leach \\& King \\shortcite{wyliki} show how to calculate the appropriate coefficient $K$ for the latter case and we plan to repeat our calculations in the future, using this prescription. In the DBA approach one assumes that the fluid constituting the accretion flow, i.e. the disc and its surroundings, is composed of blobs immersed in a dilute interblob plasma (see Aly \\& Kuijpers 1990). The blobs behave diamagnetically in the presence of the stellar magnetic field and thus suffer a surface drag force acting on them. The model has since its introduction been applied to various systems, notably to the intermediate polar class of cataclysmic variables (CV). Wynn \\& King \\shortcite{wyki} and Wynn, King \\& Horne \\shortcite{wykiho} incorporated the diamagnetic drag force into a particle hydrodynamical numerical code (HYDISC), originally developed by Whitehurst \\shortcite{wh} to simulate the accretion disc in non-magnetic close binary systems. They found important properties of the white dwarf's spin evolution and its effect on the disc and explicitly applied their findings to the moderately magnetic CV AE Aqr. A more recent study using this approach was its application to the system EX Hya \\cite{kw}. The DBA approach was extended to the study of accreting T Tauri stars for the first time by King \\& Regev \\shortcite{kr}, hereafter KR. They performed calculations of {\\em individual} blob (that is, ballistic) orbits, including the interaction with magnetic loops of the central star and have shown that this mechanism can eject blobs from the system, in directions pointing away from the disc plane. By estimating the angular momentum thus expelled and including it in the overall angular momentum budget KR found that a stellar spin equilibrium value is compatible with the observations (i.e. $\\sim$ an order of magnitude less than the breakup value) if the magnetic loops extend to a few stellar radii. Pearson \\& King \\shortcite{peki}, hereafter PK, generalised the above work into a $N$-particle simulation, in a similar manner to the one mentioned above for CVs. This work also focused on the issue of the slow observed T Tauri spin rates and confirmed the idea that such equilibrium spin rates can be achieved by the expulsion of matter from the disc until the corotation radius (see below) coincides with the edge of the magnetic loop. The additional new finding, not anticipated by KR, of this work was that the ejection of material comes to a halt as the star approaches its spin equilibrium value. The problem of the slow spin rates and rotational evolution of T Tauri stars as a result of the star's magnetic interaction with its accretion disc has been also approached, in a number of works, by analytical and semi-analytical methods, that is, stopping short of multi-dimensional numerical simulations. K\\\"onigl (1991), Cameron \\& Campbell (1993) and Armitage \\& Clarke (1996) used quite different approaches and all found that a stellar dipole magnetic field of strength $\\la 1~ {\\rm kG}$ is able to regulate the stellar spin to a quasi-static value, in the observed range. In the process of the magnetic interaction the inner accretion disc gets disrupted and the inner (abrupt) termination radius of the disc was estimated in the above papers. KR and PK used the DBA approach focusing on the ejection of mass from the disc, in an effort to link the low spin rates of T Tauri stars with outflows from young stellar objects (YSO) within a unified scenario. In the present work we utilise the DBA model in the calculation of the properties of the steady (or quasi-steady) accretion disc itself, that is, the material that remains close to the equatorial plane while spiraling in due to viscous torques, {\\em after} the spin period has already stabilised. Some pecularities of T Tauri {\\em light-curves} (i.e. the temporal luminosity variations) have already been treated using a model based on single blob ballistic orbits, which leave the the disc plane but ultimately return and impinge on the disc surface \\cite{ulregber}. In order to investigate whether the DBA model is consistent with some of the basic properties of T Tauri {\\em spectra} (see Bertout 1989 and references therein) and in particular the above mentioned IR excess, we have performed numerical simulations using a code based on HYDISC, suitably modified so as to adapt it for the problem at hand. This paper is organised as follows. In the next section we discuss the basics of blob dynamics and estimate the relevant time-scales. \\S 3 describes the numerical code used in the accretion flow simulation and the procedure for finding the spectrum emitted by such flows. In \\S 4 the results of our simulations for a score of parameter values are described in some detail and finally, \\S 5 summarises this work in comparison to the results of other approaches. ", + "conclusions": "The most prominent feature of our simulation results is the formation of a low density region (a hole) in the inner part of the disc as a result of the magnetic interaction. In addition, the density distribution becomes more spread out, as the slope of $\\log \\Sigma(R)$ decreases with increasing $k$. This behaviour is typical and essentially qualitatively independent of all the other parameters. Although the \"hole\" is the more noticeable feature (see e.g. the density maps in Fig.~4), it is less significant to overall observable features than the global change in the density distribution. We have also seen, by comparing our calculations with what can be expected from analytical results, that we uncover some features (like the absence of hot material close to the inner hole) which can not be obtained from just classical disc models with a hole inside. Due to computer power restrictions we were unable to achieve a high enough resolution, so as to see the accretion flow along the field lines and outflows from the system. In addition we were able to simulate only a limited portion of the disc and thus the calculations of the spectrum reveal only the general trend. This trend, resulting from the flattening of the surface density distribution, was always to shift emission power toward longer wavelengths. The flattening of $\\lambda F_{\\lambda}$ is most apparent near its maximum and it is reasonable that were it not for the close (computational) cutoff in $R$, the shape would remain flat to longer wavelengths. We propose therefore that the IR excess in T Tauri stars can be attributed to magnetic interaction, which modifies the functional dependence of the surface density in the surrounding discs. Existing models attempting to explain the IR excess in T Tauri stars fall into two distinct categories: those invoking geometrical factors and others, proposing energy dissipation mechanisms operating preferably in the outer parts of the disc. Our model suggests a correlation between the spectrum flatness and the magnetic field strength, appears to be quite robust (practically any shape of the magnetic field would do) and does not require assumptions about flared shapes of discs or unusual energy dissipation modes. At this stage, the results of our calculation provide little more than support for a {\\em qualitative} promising idea. As it was mentioned in the Introduction, it is possible to apply a similar prescription to the case in which the magnetic field penetrates the disc. We can reasonably expect that the results of such a calculation should not be too different, at least qualitatively, from the ones presented in this paper. It appears that all that is required is magnetic field lines imparting a torque on the gas, which changes in sign as we cross the corotation radius. Significant results for both models of the magnetic interaction can be achieved in high resolution (significantly larger number of particles). However, to extend the idea into a reliable quantitative model, full multidimensional MHD simulations, including radiative transfer, have to be ultimately performed." + }, + "0112/astro-ph0112221_arXiv.txt": { + "abstract": "The dark energy that appears to produce the accelerating expansion of the universe can be characterized by an equation of state $p=w\\rho$ with $w<-1/3$. A number of observational tests have been proposed to study the value or redshift dependence of $w$, including SN Ia distances, the Sunyaev-Zel'dovich effect, cluster abundances, strong and weak gravitational lensing, galaxy and quasar clustering, galaxy ages, the \\lya\\ forest, and cosmic microwave background anisotropies. The proposed observational tests based on these phenomena measure either the distance-redshift relation $d(z)$, the Hubble parameter $H(z)$, the age of the universe $t(z)$, the linear growth factor $D_1(z)$, or some combination of these quantities. We compute the evolution of these four observables, and of the combination $H(z)d(z)$ that enters the Alcock-Paczyznski anisotropy test, in models with constant $w$, in quintessence models with some simple forms of the potential $V(\\phi)$, and in toy models that allow more radical time variations of $w$. Measurement of any of these quantities to precision of a few percent is generally sufficient to discriminate between $w=-1$ and $w=-2/3$. However, the time-dependence predicted in quintessence models is extremely difficult to discern because the quintessence component is dynamically unimportant at the redshifts where $w$ departs substantially from its low-$z$ value. Even for the toy models that allow substantial changes in $w$ at low redshift, there is always a constant-$w$ model that produces very similar evolution of all of the observables simultaneously. We conclude that measurement of the effective equation of state of the dark energy may be achieved by several independent routes in the next few years, but that detecting time-variation in this equation of state will prove very difficult except in specialized cases. ", + "introduction": "\\label{sec:intro} The big cosmological surprise of recent years is that the dominant form of energy in the universe has negative pressure and is therefore causing the expansion of the universe to accelerate. The most direct evidence for acceleration comes from the Hubble diagram of Type Ia supernovae (SN Ia), in particular the relative apparent brightness of SN Ia at redshifts $z\\sim 0$ and $z\\sim 0.5-1$ \\citep{riess98,perlmutter99}. However, other strong arguments for a ``dark energy'' component follow from combining the cosmic microwave background (CMB) evidence for a spatially flat universe \\citep{netterfield01,pryke01} with either a minimum age $t_0 \\sim 13\\;$Gyr \\citep{vandenberg96} or dynamical evidence that the density of clustered matter is well below the critical density (see \\citealt{bahcall97,carlberg97,weinberg99b} for examples of three distinct routes to this conclusion, though there are many others). The first combination, together with a Hubble constant $H_0 \\approx 70\\;\\hubunits = (14\\;{\\rm Gyr})^{-1}$ \\citep{freedman01}, requires a component whose gravitational acceleration roughly cancels the gravitational deceleration caused by the pressureless matter, so that $t_0 \\approx H_0^{-1}$. The second combination requires that the dominant form of energy be unclustered, though it implies nothing more specific about its equation of state. A more model-dependent argument for a negative pressure component comes from the success of inflationary models with cold dark matter (CDM) and a cosmological constant ($\\Lambda$) in matching a variety of constraints from CMB anisotropies and large scale structure measurements (see \\citealt{wang01x} for a recent review). In this paper, we explore the prospects for determining the equation of state of the dark energy component through a variety of observational methods. A true cosmological constant can be treated as a vacuum energy with time-independent density and pressure related by $p=-\\rho$. Current observations favor an equation of state fairly close to this prediction \\citep{garnavich98}. However, a number of authors have considered the more general possibility that the negative pressure component is a scalar field (a.k.a.\\ ``quintessence'') with energy density determined by its potential and effective equation of state $p=w\\rho$, where $w$ can be constant or time-varying \\citep{ratra88,turner97,caldwell98}. Interest in models with time-varying $w$ has been spurred by arguments that certain simple potentials lead ``naturally'' to a negative pressure quintessence component that dominates the expansion at late times, independent of the initial conditions \\citep{zlatev99,steinhardt99}. Variants on this theme include fields with a non-standard kinetic term \\citep{kessence} or models with a complex scalar field \\citep{spintessence}. Further afield, there is the possibility that the negative pressure component is a network of frustrated topological defects \\citep{vilenkin85,spergel97}, or that cosmic acceleration arises from a breakdown of general relativity rather than the addition of a new energy component (\\citealt{mannheim01}; see also \\citealt{tegmark01}). The hope, thus far unrealized, is that one of these ideas will eventually provide a natural explanation of why the vacuum energy density is 120 orders-of-magnitude below the Planck scale and why it is comparable to the matter density at the present day, without having to resort to anthropic selection arguments \\citep{efstathiou95,martel98}. Any clear evidence that $w \\neq -1$, or, better still, that $w$ varies in time, would provide crucial clues towards understanding the physics of the dark energy. Through its influence on the cosmic expansion history, this component affects many observable phenomena, including CMB anisotropies, the \\lya\\ forest, strong and weak gravitational lensing, the anisotropy of quasar and galaxy clustering in redshift space, the ages of the oldest galaxies as a function of redshift, and standard-candle or standard-ruler measurements of the distance-redshift relation. This paper discusses these potential observational tests in a unified fashion. The equation of state determines the history of the energy density $\\rho_\\phi$, which, together with the densities $\\rho_m$ and $\\rho_r$ of matter and radiation, determines the evolution of the Hubble parameter $H(z)$ via the Friedmann equation. The history of $H(z)$ in turn determines the age of the universe $t(z)$, the growth factor of linear perturbations $D_1(z)$, and distance measures like the angular diameter distance $d_A(z)$ or luminosity distance $d_L(z)$, which are related to each other by cosmology-independent powers of $(1+z)$. Essentially all proposed tests of the properties of the negative pressure component amount to measurements of $H(z)$, $t(z)$, $D_1(z)$, or $d(z)$, or some combination of them, at redshifts accessible to a particular observational technique. We will investigate the dependence of these four quantities, and of the specific combination $H(z)d_A(z)$ that is constrained by the Alcock-Paczynski (\\citeyear{alcock79}; hereafter AP) anisotropy test, on the value and time history of $w$. Our paper joins, and, we hope, complements, a flood of recent papers that examine the prospects for specific tests and specific data sets in much greater detail. Since the strongest evidence for $\\Lambda$ or a quintessence component comes from SN Ia observations, and substantial improvements are likely from ground-based campaigns and possibly a dedicated satellite (SNAP; see {\\tt http://snap.lbl.gov}), many authors have examined the extent to which present or future SN Ia observations can constrain $w(z)$ \\citep{turner97,garnavich98,astier00,chiba00,hutererturner01,saini00, barger01,chevallier01,maor01,ng01,podariu01,wang01a,wang01b,weller01}. Because CMB anisotropy predictions depend most strongly on the sum of $\\rho_\\phi$ and $\\rho_m$ while SN Ia distances depend more nearly on the difference, the combination of these complementary observations yields much tighter constraints on the negative pressure component than either does alone \\citep{caldwell98,efstathiou99,baccigalupi01, corasiniti01,doran01}. The Sunyaev-Zel'dovich effect or size of radio sources offer alternative ways of measuring $d_A(z)$ \\citep{birkinshaw99,lima01}, and the volume-redshift test using galaxy counts constrains the combination $d_A^2(z)H^{-1}(z)$ \\citep{newman00,newman01}. The evolution of the galaxy cluster mass function can constrain the linear growth factor $D_1(z)$ \\citep{benabed01,doran01b,haiman01, newman01b,weller01b}, and population synthesis modeling of galaxy spectra can constrain $t(z)$ \\citep{lima00}. Jimenez \\& Loeb (2001) suggest that relative galaxy ages can be used to measure $dz/dt$, and thus $H(z)$. \\cite{hui99a} and \\cite{huterer01} have examined constraints on $w$ that can be obtained from weak lensing, while \\cite{hu} has considered lensing in combination with the CMB. \\cite{calvao01} have discussed constraints that could be obtained by applying the AP test to the 2dF quasar redshift survey of Boyle et al. (\\citeyear{boyle00}; for related discussions see \\citealt{hui99,cappi01,dalal01,mcdonald01}). Most of these papers have considered the potential observational constraints singly, or in pairs. The goals of our more abstract discussion, where we consider all of these observables together but do not focus on specific observational strategies, are twofold. First, we aim to understand what level of precision is necessary with any of these quantities to obtain useful constraints on $w$. Second, we want to know whether these different observables provide complementary information about the time-variation of $w$, breaking degeneracies that exist for a single measure by probing different aspects of the expansion history. Unfortunately, our conclusions on the latter point are pessimistic --- there are many different ways to measure $w$, but distinguishing a time-varying $w$ from a constant $w$ is likely to prove difficult. The papers by \\cite{wang00} and \\cite{tegmark01} also consider multiple observables, focusing on present constraints and future prospects, respectively. Tegmark's paper, in particular, is similar in spirit to ours, but different in the way that it frames the problem and evaluates the prospects. In the next section we discuss the various quintessence models that we examine in this paper. We discuss the observables in \\S\\ref{sec:observables}, beginning with the formulas that relate these quantities to the expansion history and proceeding to a brief account of observations that might measure these quantities in the next few years. We present our results in \\S\\ref{sec:results}, first for the quintessence models described in \\S\\ref{sec:quintessence}, then for a class of ``toy'' models designed to allow stronger time-variation of $w$ at low redshift. We summarize our conclusions in \\S\\ref{sec:conclusions}. ", + "conclusions": "\\label{sec:conclusions} Our results both confirm and extend earlier work on this subject. For any of the five observable quantities considered here --- the angular-diameter distance $d_A(z)$, the Hubble parameter $H(z)$, the age of the universe $t(z)$, the linear growth factor $D_1(z)$, or the Alcock-Paczynski parameter $H(z) d_A(z)$ --- measurement with $\\sim$10\\% precision near the observable's redshift of peak sensitivity would be sufficient to distinguish an $n=2$ ($w = -1/3$) model from a pure cosmological constant, even if $\\Om$ were known only to an accuracy of $\\pm 0.05$. Although this value of $w$ is already ruled out by the SN Ia measurements, our results suggest that other observations may soon be able to independently confirm the result. Distinguishing an $n=1$ ($w = -2/3$) model from a pure cosmological constant is much harder, demanding measurement precision of a few percent near the redshift of peak sensitivity, along with a determination of $\\Om$ to within $\\pm 0.05$. Although this level of precision is currently unavailable, it seems clearly within reach of improving SN Ia data, and it is likely to be achieved by one or more of the other observational methods discussed in \\S\\ref{sec:observables}. Thus, while SN Ia surveys may provide the first precise determination of $w$, a collection of other observations seems likely to provide confirmation (or refutation!) of the measurement within a few years. The sensitivity of the observables to the value of $n$ depends on redshift in different ways, reflecting the links between these quantities and the expansion history. The age $t(z)$ depends only on expansion at redshifts greater than $z$, so its sensitivity to $n$ decreases monotonically with increasing $z$. The linear growth factor, on the other hand, depends on clustering from redshift $z$ to redshift zero, so the sensitivity of $D_1(z)$ increases monotonically with $z$. The Hubble parameter $H(z)$ is most sensitive at $z\\sim 1-2$, when $\\Omega_\\phi$ is substantially different from its present-day value but not so small that quintessence is dynamically unimportant. The sensitivity of the angular diameter distance is fairly flat over a wide range of redshift. The sensitivity of the AP parameter is governed by competing effects of $H(z)$ and $d_A(z)$, which cancel each other at $z\\sim 3$. Because of their different connections to the expansion history, we hoped at the outset of this investigation that these observables would provide complementary information about the history of the equation of state, allowing a combination of measurements to detect a time-variation of $w$ that could not be found by any one method on its own. Unfortunately, we find that the level of complementarity is too weak to be useful in practice: models that make indistinguishable predictions for one observable generally make indistinguishable predictions for all of them. Of course, it is valuable to confirm an important result like a measurement of $w$ by independent methods, to check for systematic errors or a breakdown of the assumptions implicit in each approach. Also, different observables can provide complementary information about $\\Om$, precise knowledge of which is essential if one hopes to constrain $w$. However, once $\\Om$ is known, the constraints on the equation of state and its history will be dominated by the single highest precision measurement; adding lower precision measurements of other observables will give little additional purchase. We find, furthermore, that none of the observables holds much promise for distinguishing a quintessence model with a time-dependent equation of state from an appropriately chosen constant-$n$ model, even if one is highly optimistic about the achievable precision and assumes perfect, independent knowledge of $\\Om$. Tracker models with $V(\\phi)\\propto \\phi^{-1}$ and $V(\\phi)\\propto \\phi^{-6}$ are effectively identical to models with constant $n=1$ and $n=2$, respectively. Models with an Albrecht-Skordis potential cannot be distinguished from a pure-$\\Lambda$ model, except, perhaps, by a measurement of the growth factor at recombination from CMB anisotropy (a point we will return to shortly). The fundamental difficulty is that, in any observationally viable model, quintessence becomes dynamically important only at low redshift, so it affords little purchase for measuring redshift dependence of its equation of state. Furthermore, as Figures \\ref{Fig:BrokenModels_n=0}$-$\\ref{Fig:BrokenModels_n=2} demonstrate, even a substantial transition in $n$ at low redshifts is very difficult to detect, since the value of $n$ at $z=0$ is not known a priori, and time-variation must therefore be judged relative to the constant-$n$ model that best mimics the time-variable model. Our broken power-law models have substantial low redshift transitions by design, but there is usually a constant-$n$ model that predicts the same values of all observables to better than 1\\% at all observationally accessible redshifts. We conclude that detecting time-variation of the equation of state requires truly extraordinary precision, unless this variation occurs on a timescale much shorter than the Hubble time, which is possible but seems physically unlikely. Sub-percent precision may be achievable by some methods (SN Ia observations look to us like the best hope), but it requires controlling systematic uncertainties, especially those that are correlated among different redshift bins, very tightly. Our conclusions in this regard agree with those of \\cite{maor01}, who found that accurate measurements of the luminosity distance alone would be insufficient to determine the form of $w(z)$ for the dark matter energy component. \\cite{wang01a} and \\cite{tegmark01} showed that SN Ia measurements should be able to detect time-variation in the energy density $\\rho_\\phi(z)$, but this only means demonstrating that $n>0$ ($w>-1$); we agree that a significant departure from $n=0$ should be detectable, but detecting time-variation of $n$ is far more challenging. In a similar vein, despite fairly optimistic assumptions about the prospects for the SNAP satellite, \\cite{hutererturner01} find that error bars on the time-derivative of $w$ are quite large, and degrade considerably with uncertainty in $\\Omega_m$. \\cite{yamamoto01} suggest that the form of the dark energy equation of state might be determined by studying strong gravitational lensing systems, but their results indicate that detecting time-variation is possible only with extremely high precision measurements of the lensing systems, and then only if $\\Om$ is known precisely. The principal significance of our results, relative to these earlier papers, is that they apply to {\\it all} proposed observable tests based on the cosmic expansion history, since these tests always measure some combination of $H(z)$, $d_A(z)$, $t(z)$, or $D_1(z)$. Our investigation shows that there is one generic form of time-variation in the equation of state that might be observationally detectable. Constant-$n$ models with $n\\geq 2$ ($w \\geq -1/3$) are ruled out by current data, but a time-variable model could have $n \\geq 2$ at high redshift and a transition to low $n$ at low redshift when quintessence becomes the dominant energy component. The Albrecht-Skordis model displays just this behavior, since the quintessence roughly tracks the matter energy density ($n\\approx 3$) along the exponential part of $V(\\phi)$ but changes its equation of state (to $w\\approx -1$, $n\\approx 0$) when it reaches the potential minimum. If $n\\geq 2$ down to some fairly low redshift, then the dynamical effects of quintessence are non-negligible (though small) over a fair fraction of the post-recombination expansion history, and they slow the progress of matter clustering. The result is a slight (few percent) mismatch between the value of $D_1$ at $z=z_r\\approx 1100$ and the value expected for a constant-$n$ model that matches the low redshift data; in observational terms, the level of CMB anisotropy would be a few percent higher than anticipated. \\cite{doran01b} emphasize a similar point and discuss the relation between CMB anisotropy and $\\sigma_8$ in detail. Detecting even this type of time-variation will be very challenging, requiring a precise determination of the effective low-redshift value of $n$, precise determinations of the present-day amplitude of matter clustering and $\\Om$, and the demonstration that any excess CMB anisotropy does not arise from other sources, such as tensor fluctuations, secondary anisotropies, or contaminating foregrounds. The discovery of dark energy is an extraordinary cosmological achievement, one that could happen only in the era of ``precision cosmology.'' If the equation of state of this dark energy is substantially different from $p=-\\rho$, or if it has been different in the recent past, then that departure should be detected independently by several of the ambitious observational efforts currently planned or underway. A precise ($\\sim \\pm 0.1$) measurement of the low-redshift value of $w$ would be another extraordinary achievement, ruling out many models for the origin of dark energy and tightening the parameter space of others. However, the information provided by different observable probes of the cosmic expansion history, or by the same probe at different redshifts, is mostly redundant rather than complementary, once $\\Om$ has been determined to high precision. As a result, the next step of detecting time-variation in the cosmic equation of state is likely to prove extremely difficult. If we are lucky, then the dark energy has the kind of dynamical significance at high redshift or sudden transition at low redshift that produces an observationally accessible signature, though reading that signature will still require a combination of several cosmological measurements of unprecedented precision. If we are not so fortunate, then the observable effects of the dark energy will, for the foreseeable future, provide only two numbers with which to describe it, the current energy density and an effective low-redshift value of $w$ (or some equivalent pair of parameters). Until a physical model comes along that accounts for these two numbers in a natural way without adjustable inputs, the true nature of the dark energy component is likely to remain mysterious. \\vskip 0.1 in {\\bf Acknowledgments} A.M.L. and R.J.S. were supported in part by the DOE (DE-FG02-91ER40690). D.H.W. was supported in part by the NSF (AST-0098584). D.H.W. acknowledges the hospitality of the Institute for Advanced Study and financial support of the Ambrose Monell Foundation during the final phases of this work. We thank L. Amendola, R. Jimenez, E. Linder, and the anonymous referee for helpful comments on the manuscript. \\newpage" + }, + "0112/astro-ph0112017_arXiv.txt": { + "abstract": "We present preliminary results from a systematic spectral study of pulsars and their wind nebulae using the {\\it Chandra X-Ray Observatory.} The superb spatial resolution of {\\it Chandra} allows us to differentiate the compact object's spectrum from that of its surrounding nebulae. Specifically, for six Crab-like pulsars, we compare spectral fits of the averaged pulsar wind nebulae (PWN) emission to that of the central core using an absorbed power-law model. These results suggest an empirical relationship between the bulk averaged photon indicies for the PWNe and the pulsar cores; $\\Gamma_{PWN} = 0.8 \\times \\Gamma_{Core} + 0.8$. The photon indices of PWNe are found to fall in the range of $1.3 < \\Gamma_{PWN} < 2.3$. We propose that the morphological and spectral characteristics of the pulsars observed herein seem to indicate consistent emission mechanisms common to all young pulsars. We point out that the previous spectral results obtained for most X-ray pulsars are likely contaminated by PWN emission. ", + "introduction": "Recent observation of pulsars associated with supernova remnants obtained with the {\\it Chandra} X-ray observatory (Weisskopf, O'Dell \\& van Speybroeck 1996) are providing an unprecedentedly detailed view of pulsar wind nebulae. For the first time, emission features involving wisps, co-aligned toroidal structures, and axial jets are fully resolved in X-rays on arcsecond scales. These features, similar to those seen from the optical Crab nebula, are now found to be common to young, energetic pulsar in supernova remnants (Gotthelf 2001). Herein we present preliminary spectral analysis of several PWNe observed with {\\it Chandra}, which, collectively, suggest a fundamental observational relationship between the spectral characteristics of pulsars and their pulsar wind nebulae. Table 1 presents spectral results from a sample of {\\it Chandra} pulsars which have characteristic wind nebulae. A summary of these objects along with references and images can be found in Gotthelf (2001). All observations were obtained with the ACIS-CCD camera which is sensitive to X-rays in the 0.2--10~keV band with an energy resolution of $\\Delta E / E \\sim 0.1$ at $1$ keV. The on-axis point spread function is slightly undersampled by the CCD pixels ($0.5^{\\prime\\prime}$) allowing us to isolate the pulsar emission from that of the nebula. Except for N157B, which serendipitously fell on ACIS-I0, all data were obtained with ACIS-S3. The data were collected in nominal spectral (``FAINT'') and timing (3.24~s) mode and reduced and analyzed using the latest version of CIAO (CIAO 2.2/CALDB v.2.9). Starting with the Level 1 processed event files we corrected the event data for CTI effects (Townsley et al. 2001) and applied the standard Level 2 filtering criteria, then further rejected time intervals of anomalous background rates. \\begin{table} \\caption{Spectral Properties of Pulsars and their Wind Nebulae$^{a}$ } \\begin{tabular}{lccccc} \\hline Remnant$^{b}$ & $\\Gamma_{\\rm PWN}^{\\rm Averaged}$ & $\\Gamma_{\\rm Core}^{c}$ & $\\Gamma_{\\rm Pulsed}^{d}$ & log L$_{\\rm x\\,NS}$ & log L$_{\\rm x\\,PWN}$ \\\\ \\\\ \\tableline G11.2$-$0.3 & 1.28$\\pm$0.15 & 0.63$\\pm$0.12 & 0.60$\\pm$0.60 & 33.9 & 34.2 \\\\ Vela~XYZ & 1.50$\\pm$0.04 & 0.95$\\pm$0.24 & 0.93$\\pm$0.26 & 31.2 & 32.6 \\\\ Kes~75 & 1.88$\\pm$0.04 & 1.13$\\pm$0.11 & 1.10$\\pm$0.30 & 35.2 & 36.0 \\\\ 3C\\,58 & 1.92$\\pm$0.11 & 1.73$\\pm$0.15 & \\dots & 33.0 & \\dots \\\\ Crab~Nebula & 2.11$\\pm$0.05 & 1.63$\\pm$0.09 & 1.86$\\pm$0.07 & 35.9 & 37.3 \\\\ N157B~Nebula & 2.28$\\pm$0.12 & 2.07$\\pm$0.21 & 1.60$\\pm$0.35 & 35.9 & 36.1 \\\\ \\tableline \\end{tabular} \\end{table} \\noindent{$^a$}{\\footnotesize Ranked by increasing {\\bf averaged} PWN photon index.} \\\\ \\noindent{$^b$}{\\footnotesize Values for the following objects were taken from the literature: Crab: Willingale et al.~(2001); N157B: Wang \\& Gotthelf 1998.} \\\\ \\noindent{$^c$}{\\footnotesize The value for the Crab pulsar has been obtained from the literature. The measured photon-indices have been corrected for pile-up as discussed in the text.} \\\\ \\noindent{$^d$}{\\footnotesize Pulsed PI value references: Vela: Strickman, Harding \\& Jager (1999);G11.2-0.3: Torii et al.~(1997); Kes~75: Gotthelf et al.~(2000); Crab: Pravdo, Angelini \\& Harding (1997); N157B: Marshall et al.~(1998).} \\\\ \\bigskip For each object listed in Table 1, we extracted spectra from the PWN, pulsar core, and background, when available, or obtained spectral parameters from the literature as indicated. For the core spectra, the brightest central pixels were extracted based upon data above 4.0\\,keV, where the core emission is substantially greater than that of the surrounding nebula. The nebula itself was then extracted from the region above which the background was constant, excluding the core. Large variations in background rates and size of both the nebula and the central source amongst observations precluded the use of a standard extraction aperture. Custom spectral response matrix functions (RMFs) were provided with the CTI correction software, and ancillary response functions (ARFs) were created according to standard CIAO 2.2 procedures, using the QEU calibration files similarly provided. All spectra were grouped to a minimum of 50 counts per spectral bin. We fit the resultant background-subtracted pulsar and PWN spectra with a power-law model above $2$~keV using the latest version of {\\tt XSPEC} (v11.1). The spectral fits to the core included a convolution model to account for pileup effects on the spectra. Table 1 lists the spectral parameters obtained for each object; the values for the Crab and N157B have been obtained from the literature (Willingale et al.~(2001), Wang \\& Gotthelf 1998, respectively). No measurement of the pulsed spectrum for 3C58 is currently available. Absorption values were obtained by fitting the nebula spectra with an absorbed power-law above 0.6\\,keV, and were subsequently frozen for all future fits. The best-fit linear relationship between core and PWN photon indices is shown in Figure 1. We expect that the linear relationship seem in Fig. 1 may steepen slightly as we more fully account for pileup effects in the spectra. We note that the pileup corrected photon indices are likely subject to change due to the preliminary nature of the pileup model implementation, and suggest caution until these issues are resolved, but expect the basic result to remain unchanged. We also note that due to extreme pileup effects and the brightness of its nebulae compared to the brightness of its core, SNR~0540-69 has not been included pending a more detailed analysis of the available data. As a check to our final spectral results and to look for any systematic effect due to pile-up, we compare our measured power-law indices with those obtained from the literature for the pulsed emission for each source (Figure 2). Although they need not be the same, it is reassuring that they agree to within measurement errors. \\begin{figure} \\small \\begin{minipage}[t]{0.49\\linewidth} \\psfig{file=gotthelfe1_boston_fig1.ps,width=\\linewidth,angle=270.0} \\bigskip { {\\bf Figure 1. $\\Gamma_{\\rm PWN}^{\\rm Averaged}$\\,vs.\\,$\\Gamma_{\\rm Core}$} -- A plot of the power-law photon indices of the average pulsar wind nebulae spectra versus the pileup-corrected core photon indices for the collection of objects presented in Table 1. A dashed-line indicates the best-fit linear regression (i.e. $\\Gamma_{PWN} = 0.8 \\times \\Gamma_{Core} + 0.8$). The physical origin of this relationship is yet to be determined. } \\end{minipage}\\hfill \\begin{minipage}[t]{0.49\\linewidth} \\psfig{file=gotthelfe1_boston_fig2.ps,width=\\linewidth,angle=270.0,clip=} \\bigskip { {\\bf Figure 2. $\\Gamma_{\\rm Core}^{\\rm Pulsed}$\\,vs.\\,$\\Gamma_{\\rm Core}$} -- A comparison between the subset of measured power-law photon index for the core emission of objects presented in Table 1 and the pulsed emission for each of these source obtained from the literature (from ASCA or XTE fits). A one-to-one correspondence is indicated by the dashed line. } \\end{minipage} \\end{figure} ", + "conclusions": "" + }, + "0112/astro-ph0112367_arXiv.txt": { + "abstract": "We present the analysis of \\easca archival data from the Galactic source W51. The \\easca spectra show that the soft ($kT\\simlt 2.5$ keV) X-rays are of thermal origin and are compatible with W51C being a single, isothermal ($kT\\simeq 0.3$~keV) supernova remnant at the far-side of the Sagittarius arm. The \\easca\\ images reveal hard ($kT\\simgt 2.5$ keV) X-ray sources which were not seen in previous X-ray observations. Some of these sources are coincident with massive star-forming regions and the spectra are used to derive X-ray parameters. By comparing the X-ray absorbing column density with atomic hydrogen column density, we infer the location of star-forming regions relative to molecular clouds. There are unidentified hard X-ray sources superposed on the supernova remnant and we discuss the possibility of their association. ", + "introduction": "W51 is an extended ($\\sim 1^\\circ$) radio source located at the tangential point ($l=49^\\circ$) of the Sagittarius arm. It is composed of two complex H~II regions, W51A and W51B, and the supernova remant (SNR) W51C (e.g., Bieging 1975; Koo 1997). W51A forms the northern part of W51, and is separated from the other two sources. It contains two major components, each composed of several compact H~II reigons. W51B is composed of at least six compact H II regions scattered over an area of $\\sim 15'$ size. W51B sources are associated with a stream of atomic and molecular gases, with line-of-sight velocities significantly greater than the maximum velocity permitted by Galactic rotation alone. The stream is thought to be gas flowing along the Sagittarius spiral arm in response to the perturbation due to the spiral potential (e.g., Burton 1971). Superimposed on W51B, there is an extended structure W51C which is a SNR. W51C appears in radio continuum as an incomplete shell of $\\sim 30'$ extent with its upper portion open \\citep{cop91, sub95}. Shocked atomic and molecular gases have been detected in the western part of the SNR, which indicates that the SNR is interacting with a large molecular cloud \\citep{koo91, koo97a, koo97b}. The complex structure of W51 is partly due to the inclination of the Sagittarius arm, so that we look down the length of the arm a distance of 5~kpc along the line-of-sight. Soft ($\\simlt 2$ keV) X-ray emission associated with W51 has been detected by {\\em Eintesin} and ROSAT \\citep[hereafter KKS]{sew90, koo95}. Diffuse X-rays come from a region surrounding W51B and W51C, whereas W51A has essentially no soft X-ray emission associated with it. The soft X-ray emitting region is elongated ($50'\\times 38'$) along the east-west direction, and may be divided into three parts; a central structure composed of two bright regions separated by $\\sim 10'$, an incomplete shell $\\sim 30'$ long in the east, and an extended ($\\sim 20'\\times 10'$) structure in the west. The structure in the west is separated from the central region by a region of weak emission where molecular clouds associated with the W51B star-forming region are located. The systematic hardening of the X-ray spectrum toward the west suggested that the X-rays are emitted behind these molecular clouds. The average X-ray spectrum of the SNR was fitted well by a single-temperature ($k_B T\\simeq 0.29$~keV) thermal plasma model. But the energy range of the ROSAT detector was not enough to distinguish this from other emission models. In this paper, we present the results of an \\easca study of the W51 complex. The greater (0.7--10 keV) energy coverage of \\easca makes it clear that the X-ray emission from the SNR is of thermal origin. The newly derived plasma parameters are consistent with the ROSAT results. \\easca also reveals hard X-rays from star-forming regions in W51, which were not seen by ROSAT. ", + "conclusions": "W51 is a complex region with a SNR and star-forming regions superposed within an area of angular diameter $\\sim 1^\\circ$. Soft X-rays from the SNR were detected by {\\em Einstein} and ROSAT, but, because of the limited energy range of detector, it was not conclusive whether the emission was thermal or non-thermal. The \\easca study in this paper clearly shows that most parts of the SNR emit thermal X-rays and the temperature of the hot gas is about $0.3$~keV. It also shows that the X-rays originate behind most of the atomic and molecular gas along those line of sights. These results are compatible with X-rays being emitted from a single, isothermal SNR at the far-side of the Sagittarius arm, which makes the W51C SNR one of the largest SNRs in the Galaxy. The \\easca observations reveal hard X-rays from the W51B star-forming regions. Two regions (G48.9$-$0.3 and G49.0$-$0.3) appear bright in hard X-rays and it is likely that they are located in the front side of the molecular cloud. W51 is a rare example of a SNR superposed on a star-forming region. Since the progenitors of Type II SN are massive stars, this situation should be more common, but it is not. Apparently W51 is one of those few regions in the Galaxy where we can study the formation of massive stars and the consequence of their violent explosions simultaneously. High-resolution X-ray observations are needed to fully identify the origins of the central emission." + }, + "0112/astro-ph0112198_arXiv.txt": { + "abstract": "{ Most of the current theories suggest the \\LB\\ phenomenon to originate from an interaction between the stellar surface and its local environment. In this paper, we compare the abundance pattern of the \\LB\\ stars to that of the interstellar medium and find larger deficiencies for Mg, Si, Mn and Zn than in the interstellar medium. A comparison with metal poor post-AGB stars showing evidence for circumstellar material indicates a similar physical process possibly being at work for some of the \\LB\\ stars, but not for all of them. Despite the fact that the number of spectroscopically analysed \\LB\\ stars has considerably increased in the past, a test of predicted effects with observations shows current abundance and temperature data to be still controversial. ", + "introduction": "In the last few years, the \\LB\\ stars (metal-poor population~{\\sc i} A to F type stars) have experienced increased attention by abundance analysis groups. The results have been collected by \\citet[ hereafter referred to as Paper~I]{Heit:01a} and show that the proportion of \\LB\\ stars with known abundances is now large enough to examine the abundances with respect to other stellar parameters on a good statistical basis. The analysed stars span a wide range of atmospheric parameters, in particular the effective temperature (Fig.~\\ref{evol}). This parameter plays a major role in current theories, which are briefly reviewed in the following. \\citet{Venn:90} proposed that the peculiar abundance patterns of \\LB\\ stars originate from the interstellar medium, which shows a similar abundance distribution (see Sect.~\\ref{ISM}). Within this hypothesis it is assumed that only the interstellar gas, but not the dust, is accreted onto the surface of the stars. \\citet{Char:91a} calculated the concentration of the elements Ca, Ti, Mn and Eu in the superficial convection zone (SCZ) within a simple analytical model, which takes into account accretion of interstellar gas and diffusion below the SCZ, for various effective temperatures. It is assumed that the atmospheric material is mixed thoroughly from the surface to the bottom of the SCZ. More detailed numerical calculations have been performed by \\citet{Turc:93} for the elements Ca, Sc and Ti. Abundance profiles of the first two elements show overabundances at the surface if only chemical separation and convective mixing is taken into account, whereas Ti is predicted to be underabundant in this case. If accretion of circumstellar gas with a certain amount of depletion is added, the calculations show that for an accretion rate $\\dot M$ of at least $5\\cdot 10^{-14}$~M$_{\\odot}$~yr$^{-1}$, the abundances of the examined elements in the convection zone converge to the values in the accreted gas on a very short timescale. Two other points became evident. Independently of $\\dot M$ and the duration of the accretion phase, the abundances are again governed by chemical separation when accretion is stopped. This means that accretion must be an ongoing process if it is responsible for the observed abundance pattern. Secondly, meridional circulation induced by rotation with equatorial velocities up to 125~\\kms\\ does not alter the surface abundances produced by accretion. Higher rotation rates could not be treated due to numerical problems. All these calculations are based on only one static stellar envelope model with a fixed parameter set of (\\Teff, \\logg) = (8000~K, 4.3). The separation of gas and dust in a circumstellar shell was investigated by \\citet{Andr:00}, who calculated gas and dust grain velocities in a shell extending to 100 stellar radii around a star with \\Teff=8500~K, assuming a polytropic density distribution. For a ratio between radiative and gravitational acceleration on the gas of 0.99, large dust grains and a rather smooth density distribution (polytropic index =~2), they indeed find dust grains to be forced to an outward motion by radiative pressure. The separation becomes effective at a distance from the stellar surface where the temperature is about 1600~K (condensation temperature for heavy elements), which corresponds to about 10 stellar radii. The gaseous part of the shell is accreted to the surface of the star. Thus the two components are decoupled and the superficial chemical composition is changed according to the depletions in the gas coming from the outer part of the shell. The calculations take into account only interactions between neutral particles because they are shown to be more important than Coulomb-type interactions. Within this simple model only rough estimates for the gas-dust separation can be made, which are based on very restricted assumptions. But it could serve for more sophisticated models, which in particular should be extended to lower temperatures and smaller dust grains, which are more likely to be formed around \\LB\\ stars. An earlier theoretical approach to the \\LB\\ phenomenon had its origin in recalling the physical processes operating in the atmospheres of Am stars. The Am abundance pattern has been explained by \\citet{Char:91b} by chemical separation of elements below the superficial hydrogen convection zone, caused by diffusion processes. In order to produce the actual abundance values of Am stars, an additional process is needed, e.g. a small amount of mass-loss ($10^{-15}$~M$_{\\odot}$~yr$^{-1}$). Evidence for mass-loss has not yet been observed for Am stars. By introducing a two orders of magnitude higher mass-loss rate, \\citet{Mich:86} have changed the calculated Am abundance pattern to a \\LB\\ like one, but they have not been able to produce underabundances as low as were observed in several \\LB\\ stars. The underabundances even vanish for most elements if meridional circulation induced by high rotational velocities is taken into account \\citep{Char:93}. Although large uncertainties are still involved in the modeling (above all for the radiative acceleration), this theory has been widely discarded as an explanation of the \\LB\\ star abundances. Further theoretical considerations include \\citet{Andr:97}, who proposed that \\LB\\ stars are the result of a merger of contact binaries of W~UMa type. He argues that mass loss during the merger phase could form the circumstellar shell, whose accretion leads to the observed underabundances. The hypothesis is substantiated by lifetime and number estimates. \\citet{Fara:99} suspect that a part of the \\LB\\ stars are undetected spectroscopic binary systems, and that their abundance anomalies are due to veiling effects in the composite spectra. Summarizing, in the mentioned theories the \\LB\\ phenomenon seems to originate from an interaction between the stellar surface and its local environment. In the following we confront predicted effects with observations. ", + "conclusions": "\\begin{table} \\caption{Summary of the search for signatures of accretion. $\\sigma$ \\dots\\ cool stars with small (+) or large ($-$) $\\sigma$ from Table~\\ref{par_all_tab}, hot stars with [Ca]=[Ti] ({\\sf x}); I \\dots\\ all abundances higher than in ISM (+), some abundances lower than in ISM ($-$); C \\dots\\ smooth (+) or discontinuous ($-$) relation of abundances and $T_c$.} \\label{comb} \\begin{tabular}{rccc|rccc} \\hline\\hline HD & $\\sigma$ & I & C & HD & $\\sigma$ & I & C \\\\ \\hline 319 & + & & + & 125162 & $-$ & & $-$ \\\\ 11413 & $-$ & + & $-$ & 142703 & + & $-$ & + \\\\ 15165 & $-$ & & $-$ & 168740 & + & & \\\\ 31295 &{\\sf x}& + & + & 170680 & & + & \\\\ 74873 &{\\sf x}& & $-$ & 183324 &{\\sf x}& & $-$ \\\\ 75654 & & & + & 192640 & + & & $-$ \\\\ 84123 & + & $-$ & $-$ & 193256 & $-$ & & $-$ \\\\ 84948A/B & + & $-$ & & 198160/1 & $-$ & & \\\\ 101108 & + & + & $-$ & 204041 & $-$ & + & + \\\\ 106223 & + & & + & 210111 & $-$ & & + \\\\ 107233 & + & & & 221756 &{\\sf x}& & \\\\ 110411 &{\\sf x}& & $-$ & & & & \\\\ \\hline\\hline \\end{tabular} \\end{table} The \\LB\\ star abundances were examined with regard to correlations to the stellar parameters of this group, in particular the effective temperature. It was found that for some elements (C, Na, Mg, Si, Ca, Sc, Cr, Fe, Sr) the abundances are weakly correlated with \\Teff, \\logg, the age, the pulsational period or \\Vsini. The scatter of heavy element abundances in individual stars does not depend on \\Teff. These findings are inconclusive with regards to testing the accretion/diffusion theory. Because of the lack of calculations for more than three elements and different atmospheric parameters, the uncertainties related to the treatment of convection and the calculation of the radiative acceleration, and the free parameters (mainly the accretion rate and the abundance spectrum in the accreted material), we consider the theoretical models to be rather simple and incomplete. The chemical composition of the \\LB\\ stars has been compared to that of the interstellar medium (ISM). The {\\em mean} abundances of some elements (Mg, Si, Mn, Zn) are slightly lower in the \\LB\\ stars than in the ISM (by 0.2 to 0.6~dex), and the {\\em lowest} abundances found in \\LB\\ stars for these elements are lower than the lowest ISM abundances by 0.4 to 0.8~dex. Similar deviations have been found for only half of the single stars which can be compared to nearby sight lines. Within an accretion/diffusion scenario, the abundances of the accreted elements would be expected to be greater than in the ISM. The \\LB\\ abundance pattern has also been compared to that of stars with circumstellar material (post-AGB, Vega-like and A-shell stars). Similar relations of abundances with condensation temperatures suggest that the same physical processes lead to the chemical compositions of some \\LB\\ stars and the metal poor post-AGB stars although theoretical calculations for the latter group do not exist. More observations are clearly needed to confirm this hypothesis. On the other hand, the lack of metal deficiency in dusty stars with atmospheric parameters similar to \\LB\\ stars questions the connection of circumstellar dust with the \\LB\\ phenomenon, although the comparison is based on a very small sample of Vega-like stars. From the currently available abundance data we conclude that the stars HD\\,319, HD\\,31295 and HD\\,106223 could well have experienced accretion of circumstellar gas (see Table~\\ref{comb}), which however has not been detected in their spectra. For the other stars, further examination and spectral data are required." + }, + "0112/astro-ph0112151_arXiv.txt": { + "abstract": "{We examine the optical properties of the nuclei of low luminosity radio-galaxies using snapshot HST images of the B2 sample. In agreement with the results obtained from the analysis of the brighter 3C/FR~I sample, we find a correlation between fluxes (and luminosities) of the optical and radio cores. This provides further support for the interpretation that the optical nuclear emission in FR~I is dominated by synchrotron emission and that accretion in these sources takes place in a low efficiency radiative regime. In the framework of the FR~I/BL Lacs unified scheme, we find that the luminosity difference between FR~I and BL Lac nuclei can be reproduced with a common beaming factor in both the radio and the optical band, independent of the extended radio luminosity, thus supporting such a scenario. The corresponding bulk Lorentz factor is significantly smaller than is expected from observational and theoretical considerations in BL Lacs: this can be interpreted as due to a velocity structure in the jet, with a fast spine surrounded by a slower layer. ", + "introduction": "\\label{intro} The presence of a radio-source represents a common manifestation of nuclear activity associated to elliptical galaxies, in particular for the brightest members of this class. For example, among galaxies brighter than $M_{\\rm B} < -21$, more than 20 \\% have radio luminosities $L_{\\rm 408 MHz} > 10^{23.5}$ W Hz$^{-1}$ (Colla et al. \\cite{colla}). Due to the steepness of the radio luminosity function, most of them are low luminosity radio-sources and show the characteristic edge-darkened FR~I radio morphology (Fanaroff and Riley \\cite{fr}). However, in observing bands other than the radio, emission related to their active nuclei remains largely elusive. Optical spectra of low luminosity radio-galaxies are in fact dominated by the stellar component of the host galaxy with only faint narrow emission lines, while permitted broad lines are seldom detected. Similarly, most of their X-ray emission originates in the ambient thermal plasma. With the limited information available, we cannot effectively constrain the physical properties and the emission processes at work in these AGNs, and to determine how these sources fit into the AGN unification schemes: one has to rely on the properties of the host galaxies, of the environment or of the extended radio structure. Significant progress in our understanding of low luminosity radio-galaxies has been achieved recently thanks to high resolution HST and Chandra images, which enabled us to isolate their genuine optical and X-ray nuclear emission. In HST images unresolved nuclear sources are detected in the great majority of the 33 FR~I belonging to the 3CR sample (Chiaberge et al. \\cite{chiaberge99}, hereafter CCC99). The optical flux density of these Central Compact Cores (CCC) shows a striking linear correlation with the radio core, arguing for a common non--thermal synchrotron origin. In five sources, for which it is possible to estimate the viewing angle based on the inclination of their circumnuclear dusty disks, the luminosity of the central source shows a suggestive dependence on the radio galaxy orientation, as qualitatively expected if the optical emission is indeed produced in a relativistic jet (Capetti and Celotti \\cite{ac2}). The high rate of CCC detection suggests that a geometrically thick torus can be present at most in a minority of low luminosity radio galaxies. CCC fluxes also represent upper limits to any thermal/disk emission that translate (for a $10^9 M_{\\sun}$ black hole) into a fraction $\\sim 10^{-7}-10^{-5}$ of the Eddington luminosity, suggesting that accretion takes place at low rate or in a low efficiency radiative regime. This information also offers a new possibility of testing the FR~I/BL Lac unification scheme, by directly comparing the optical nuclear properties of radio galaxies with their putative aligned (beamed) counterparts, analogously to the procedure followed in the radio band (Kollgaard et al. \\cite{koll96}). From this comparison Chiaberge et al. (\\cite{chiaberge00b}, hereafter CCCG00) found that the difference in luminosity (in radio and optical bands) between FR~I nuclei and BL Lac is significantly smaller than would be expected, in the frame of a simple one--zone model, from the high bulk Lorentz factor of BL Lacs jets implied by observational and theoretical considerations (e.g. Dondi \\& Ghisellini \\cite{dondi}, Ghisellini et al. \\cite{gg98}, Tavecchio et al. \\cite{taold}). In order to reconcile these results with the unification scheme, they suggested that a velocity structure is present in the jet, with a fast spine surrounded by a slow layer, as already suggested by other evidence on larger scales (e.g. Laing 1993, Laing et al. \\cite{laing99}). We recently obtained HST images for more than half of the B2 sample of low luminosity radio galaxies (see Capetti et al. \\cite{capetti00}, hereafter Paper I). Therefore, it is now possible to complement the analysis performed for the 3C sources with a study of sources at lower radio luminosities. This extension will allow us to test the general validity of these results for the radio galaxy population and to explore in more detail the relationship between FR~I and BL Lacs. Thus in this paper we focus on the properties of their optical nuclei. In Sect. \\ref{thesample} we briefly present the properties of the B2 sample and the HST observations on which this study is based; the optical nuclear properties of the galaxies are described in Sect. \\ref{CCC} where we also quantify the contribution of their central optical sources. In Sect. \\ref{discussion} we discuss the implications of our results, which are summarized in Sect. \\ref{summary}. For consistency with Paper I we use H$_{\\rm o}$= 100 km s$^{-1}$ Mpc$^{-1}$ and $q_0=0.5$. ", + "conclusions": "\\label{summary} We discuss the optical properties of the nuclei of low luminosity radio galaxies as derived from the snapshot HST images of the B2 sample. A nuclear optical component is found in 18 out of the 57 observed galaxies and, in agreement with the results obtained from the brighter 3C/FR~I sources, we found a correlation between fluxes (and luminosities) of the optical and radio cores. This provides further support to the interpretation of a synchrotron origin of the optical nuclear emission. In the sources in which we failed to detect an optical core, the optical limits are consistent with the interpretation that their nuclei follow the same radio/optical correlation, but their nuclear flux density is insufficient to be seen against the galaxy background. However, the large fraction of undetected CCC does not allow us to strengthen our conclusions on the lack of obscuring nuclear matter (as it was in the case of the more powerful 3CR FR~I). The radio/optical nuclear correlation for all FR~I extends down to an optical luminosity of 10$^{18}$ W Hz$^{-1}$. This value can be adopted as an upper limit to any emission from the AGN not directly related to their relativistic jets, such as e.g. from the accretion disk. For a 10$^9$ M\\sun\\ central black hole, this corresponds to a fraction $\\sim 3\\times 10^{-8}$ of the Eddington luminosity in the optical band and it provides a clear evidence that accretion in these low luminosity radio galaxies occurs at a very low rate and/or radiation efficiency. In the framework of the FR~I/BL Lacs unified scheme, the direct comparison of the optical nuclear properties of radio galaxies with their putative aligned counterparts provides a quantitative method of testing this unification model. The inclusion of the B2 sources significantly improves the coverage towards low extended luminosities with respect to the 3C/FR~I sample and thus it is now possible to compare the properties of the two populations over their whole range of radio power. This comparison indicates that the differences in luminosity between BL Lac and FR~I can be explained with a single amplification factor over the whole range of extended radio luminosity in both the radio and optical bands, again supporting the interpretation that in both cases we are seeing relativistic jet synchrotron emission. The corresponding bulk Lorentz factor results (for typical viewing angles) significantly smaller than derived from spectral energy properties. This support the interpretation that a velocity structure is present in the jet, with a fast spine surrounded by a slower layer where the layer is responsible for the bulk of the intrinsic emission." + }, + "0112/hep-ph0112134_arXiv.txt": { + "abstract": "In the framework of the Constrained Minimal Supersymmetric Standard Model we discuss the impact of the recent experimental information, especially from the E821 Brookhaven experiment on $g_{\\mu}-2$ along with the light Higgs boson mass bound from LEP, in delineating regions of the parameters which are consistent with cosmological data. The effect of these to the Dark Matter direct searches is also discussed. ", + "introduction": "Supersymmetry (SUSY) is a landmark in our efforts to construct a unified theory of all fundamental interactions observed in nature. At very high energies, close to the Planck scale ($M_P$) it is indispensable in constructing consistent string theories, and at low energies ($\\sim 1 \\TeV$) it seems unavoidable if the gauge hierarchy problem is to be resolved. Such a resolution provides a measure of the supersymmetry breaking scale $M_{SUSY} \\thickapprox \\mathcal{O}(1 \\TeV) $. There is indirect evidence for such a low-energy supersymmetry breaking scale, from the unification of the gauge couplings \\cite{Kelley} and from lightness of the Higgs boson as determined from precise electroweak measurements, mainly at LEP \\cite{EW}. Furthermore, such a low energy SUSY breaking scale is also favored cosmologically. As is well known, $R$-parity conserving SUSY models, contain in the sparticle spectrum a stable, neutral particle, identifiable with the lightest neutralino ($\\lsp$), referred to as the LSP \\cite{Hagelin}. It is important \\cite{Hagelin} that such a LSP with mass, as low-energy SUSY entails, in the $ 100\\GeV - 1\\TeV $ region, may indeed provide the right form and amount of the highly desirable astrophysically and cosmologically Dark Matter (DM). The latest data about Cosmic Microwave Background (CMB) radiation anisotropies \\cite{cmb} not only favour a flat ($k=0$ or $\\Omega_0=1$), inflationary Universe, but they also determine a matter density $\\Omega_M h_0^2 \\thickapprox 0.15 \\pm 0.05$. Taking into account the simultaneously determined baryon density $\\Omega_B h_0^2 \\thickapprox 0.02$, and the rather tiny neutrino density, they result to \\beq \\Omega_{DM} h_0^2 = 0.13 \\pm 0.05 \\,. \\label{cbound} \\eeq If we assume that all DM is supersymmetric due to LSP, i.e. $\\Omega_{DM} \\equiv \\Omega_{\\lsp}$, it is tempting to combine the bound of Eq.\\ref{cbound} with other presently available constraints from particle physics, such as the lower bound on the mass of the Higgs bosons ($m_{h} \\geq 113.5 \\GeV$) provided by LEP \\cite{LEP} and the recent results from the BNL E821 experiment \\cite{E821} on the anomalous magnetic moment of the muon ($\\delta \\alpha_{\\mu} =43 (16)\\times 10^{-10}$). Although the situation regarding the $g_\\mu -2$ has not been definitely settled, supersymmetry emerges as a prominent candidate in explaining the discrepancy between the Standard Model predictions and experimental measurements, and in the sequel we concede that this deviation accounted for SUSY. We find that this combination of the experimental information from high energy physics and cosmology puts austere bounds on the parameter space of the Constrained Minimal Supersmmetric Standard Model (CMSSM), enabling us to investigate the potential of discovering SUSY, if it is based on CMSSM, at future colliders and direct DM search experiments. ", + "conclusions": "Concluding, we combined recent high energy physics experimental information, like the anomalous magnetic moment of the muon measured at E821 Brookhaven experiment and the light Higgs boson mass bound from LEP, with the cosmological data for DM. By doing so we studied the imposed constraints on the parameter space of the CMSSM and hence we assessed the potential of discovering SUSY, if it is based on CMSSM, at future colliders and DM direct searches experiments. The bounds put on the sparticle spectrum can guarantee that in LHC but also in a $e^{+}e^{-}$ linear collider with center of mass energy $\\sqrt{s} = 800\\GeV$, such as TESLA, CMSSM can be discovered. The guarantee for a linear collider with this energy is lost in a charged sparticle final state channel, if the lower bound on the value of $g_\\mu -2 $ is lowered to its $\\approx 2 \\sigma$ value, but not for the LHC. In this case only by increasing the center of mass energy to be $\\simeq 1.2 \\TeV$, a $e^{+}e^{-}$ linear collider can find CMSSM in $\\tilde{\\tau} \\, {\\tilde{\\tau}}^*$ or ${\\tilde \\chi}^{+} {{\\tilde \\chi}^{-}} $ channels. The impact of the E821 experiment's result along with the bound on Higgs mass is also significant for the direct DM searches. We found that the maximum value of the spin-independent $\\lsp$-nucleon cross section attained is of the order of $10^{-8}$ pb. Moreover this cross section can not be lower than $10^{-10}$ pb, which is very promising for the forthcoming direct DM experiments." + }, + "0112/astro-ph0112401_arXiv.txt": { + "abstract": "We present results of a fully non-local, compressible model of convection for A-star envelopes. This model quite naturally reproduces a variety of results from observations and numerical simulations which local models based on a mixing length do not. Our principal results, which are for models with \\teff\\ between 7200~K and 8500~K, are the following: First, the photospheric velocities and filling factors are in qualitative agreement with those derived from observations of line profiles of A-type stars. Second, the \\heii\\ and \\hi\\ convection zones are separated in terms of convective flux and thermal interaction, but joined in terms of the convective velocity field, in agreement with numerical simulations. In addition, we attempt to quantify the amount of overshooting in our models at the base of the \\heii\\ convection zone. ", + "introduction": "Over the last five decades the most frequently used approach to describe stellar convection has been the mixing length theory (MLT, Biermann 1948, B\\\"ohm-Vitense 1958). However, the great simplicity achieved by describing convection in terms of local variables is only attained at the cost of trade-offs, the most important of which is the specification of a mixing length that can neither be derived from rigorous theory nor from observations. More recently, turbulence models, e.g.\\ by Canuto et al.\\ (1996, hereafter the CGM model) have been used to improve the MLT expressions. These convection models still provide a local expression for the temperature gradient and contain the specification of a scale length $l$. The latter also holds for non-local versions of the MLT which were proposed to account for convective overshooting. However, the intrinsic non-locality of this problem has prohibited a satisfactory solution within the context of models that use any form of local scale length (see Renzini 1987 and Canuto 1993). This difficulty is naturally avoided by numerical simulations which have come into use during the last decade as a tool to study stellar surface convection. Simulations in 3D have mostly been devoted to solar convection (Nordlund \\& Dravins 1990, Atroshchenko \\& Gadun 1994, Kim \\& Chan 1998, Stein \\& Nordlund 1998), while 2D simulations have been used for more extended computations over the HR diagram (cf.\\ Freytag 1995 and Freytag et al.\\ 1996). Such calculations can include the entire convective part of a stellar envelope only for the case of A-stars (and some types of white dwarfs). Even then, the computational efforts become considerable, especially when realistic microphysics is used and thermally relaxed solutions are required. To use simulations for complete stellar models is thus beyond the range of present computer capabilities (cf.\\ Kupka 2001). Another alternative was pioneered by Xiong (1978) who used the Reynolds stress approach. This approach had previously been applied in atmospheric as well as in engineering sciences. But even in its most recent version (Xiong et al.\\ 1997) his formalism still uses a mixing length to calculate the dissipation rate $\\epsilon$ of turbulent kinetic energy. Canuto (1992, 1993) and Canuto \\& Dubovikov (1998, hereafter CD98) abandoned the use of a mixing length in their Reynolds stress models. These models provide both the mean quantities of stellar structure (temperature $T$, pressure $P$, luminosity $L$, and mass $M$ or radius $r$) as well as the second order moments (SOMs) of temperature and velocity fields created by stellar convection (turbulent kinetic energy $\\rho K$, temperature fluctuations $\\overline{\\theta^2}$, convective flux $F_{\\rm C}= c_p \\rho \\overline{w\\theta}$, vertical turbulent kinetic energy $\\frac{1}{2} \\rho \\overline{w^2}$, and the dissipation rate $\\epsilon$) as the solution of coupled, non-linear differential equations. Their models are thus fully non-local on the level of second order moments. Numerical solutions of these models for the case of idealized microphysics have been presented by Kupka (1999a) and Kupka (1999b, 2001). The same equations, using realistic microphysics, were later solved for the \\heii\\ convection zone of A-stars (Kupka \\& Montgomery 2001; these results were first discussed in Canuto 2000). In this paper, we present solutions for complete A-star envelopes. Numerically, this problem is easier than that of convection in the Sun since A-stars are hotter and therefore have thinner convection zones. In addition, A-stars reveal the shortcomings of local convection models more clearly, as their less efficient convection is much more sensitive to details in the modelling. Depending, for example, on whether an $\\alpha$ of 0.5 or 1.5 is chosen in MLT for a main sequence star with \\teff\\ $\\sim 7500$~K, an envelope may either have a mostly radiative temperature gradient or still contain a nearly adiabatic region. This holds for any of the convection models which rely on a convective scale length. Hence, the efficiency of convection in the envelopes of A-stars has remained an open problem and makes them a logical as well as a promising starting point for our study. In the following, we give an outline of the physics and the numerical procedure used to compute our envelope models (the discussion of the moment equation formalism is self-contained, so that readers unfamiliar with it can skip ahead without difficulties). Results are then presented for a sequence of models which differ from each other only in \\teff. We include a model with lower gravity in order to illustrate the effect of a change in $\\log g$. Finally, we show that the non-local convection model agrees with the known observational constraints and the results of numerical simulations, whereas local models are fundamentally unable to do this. ", + "conclusions": "\\label{Sect5} Using a fully non-local, compressible convection model together with a realistic equation of state and opacities, we have calculated envelope models for stellar parameters appropriate for A-stars. In examining the results of this model, we have found many points of agreement both with observations and with numerical simulations. First, our photospheric velocities are consistent with the lower limit of the typical micro- and macroturbulence parameters found for A-stars (1.5--2 km s$^{-1}$, see Varenne \\& Monier 1999 and Landstreet 1998). Line blanketing should further increase these values. We expect a smoother $v_{\\rm C}(r)$ (without small minima as in Figures~\\ref{models}b and \\ref{mltcomp}) from an improved treatment of fourth order moments and inclusion of $\\overline{p'w}$ (cf.\\ Sect.~\\ref{Sect2}). Second, we find that the filling factor for rising fluid elements in the photospheres of our models is less than 1/2, also in agreement with observations of line profiles in A-stars. Third, we find in the temperature range 7200~K to 8500~K that the \\heii\\ and \\hi\\ zones are well-separated in terms of the convective flux but {\\em not} in terms of the convective velocity field. The two zones are thus in some sense thermally separated but dynamically joined. This feature is also shown by the numerical simulations. Finally, we find an OV at the base of the \\heii\\ convection zone of $\\sim 0.45 H_p$. The numerical simulations find an even larger OV, but this may also be due to the fact that they were done in 2D. We note that in all cases we find a nearly radiative temperature gradient in the OV region, whereas the velocities in this region remain quite large, within an order of magnitude of their maxima within the convection zone ($\\sim 0.5$ km s$^{-1}$). In addition, the non-local model yields smaller temperature gradients than the local model of Canuto et al.\\ (CGM, 1996). Such a comparison with MLT is more difficult due to the large range of $\\alpha$ in current use. Nevertheless, we have found evidence that for main sequence models $\\alpha$ has to be decreased from values of $\\sim 1.0$ at about 7100~K to $\\sim 0.4$ for models with $\\teff\\ = 8000$~K in order to obtain a comparable value of $(F_{\\rm C})_{\\rm max}$ in the \\hi\\ convection zone. In order to match $(F_{\\rm C})_{\\rm max}$ in the \\heii\\ convection zone, a completely different set of $\\alpha$'s (with larger values) would be required. As already mentioned, A-stars are excellent choices for this first calculation since they have relatively thin surface convection zones, so that the thermal time scales involved are not so long. In addition, they are interesting stars in their own right, containing high-metallicity stars (the Am stars) as well as two groups of pulsating stars (the roAp and $\\delta$ Scuti stars). In the future, it may be possible to use the pulsating stars as probes of the subsurface convection zones, much as has been done in the case of the Sun." + }, + "0112/astro-ph0112292_arXiv.txt": { + "abstract": "The galaxy pair NGC 5194/95 (M51) is one of the closest and best known interacting systems. Despite its notoriety, however, many of its features are not well studied. Extending westward from NGC~5195 is a low surface brightness tidal tail, which can only be seen in deep broadband exposures. Our previous [O~III] $\\lambda 5007$ planetary nebulae (PN) survey of M51 recovered this tidal tail, and presented us with a opportunity to study the kinematics of a galaxy interaction in progress. We report the results of a spectroscopy survey of the PN, aimed at determining their kinematic properties. We then use these data to constrain new self-consistent numerical models of the system. ", + "introduction": " ", + "conclusions": "" + }, + "0112/astro-ph0112547_arXiv.txt": { + "abstract": "Elliptical and S0 galaxies dominate the galaxy population in nearby rich clusters such as Coma. Studies of the evolution of the colors, mass-to-light ratios, and line indices of early-type galaxies indicate that they have been a highly homogeneous, slowly evolving population over the last $\\sim 65$\\,\\% of the age of the Universe. On the other hand, recent evidence suggests that many early-type galaxies in clusters have been transformed from spiral galaxies since $z \\sim 1$. Arguably the most spectacular evidence for such transformations is the incidence of red merger systems in several high redshift clusters. Due to this morphological evolution the sample of early-type galaxies at high redshift is only a subsample of the sample of early-type galaxies at low redshift. This ``progenitor bias'' results in an overestimate of the mean formation redshift if simple models without morphological transformations are used. Models which incorporate morphological evolution explicitly can bring the homogeneity, slow evolution, and morphological transformations into agreement. The modeling shows that the corrected mean formation redshift of the stars in early-type galaxies may be as low as $z\\approx 2$ in a $\\Lambda$ dominated Universe. ", + "introduction": "The galaxy population in rich clusters is dominated by early-type galaxies (S0 and elliptical galaxies). The study of these objects gives insight in the formation of the most massive disk- and spheroidal galaxies in the Universe, and in the processes governing star formation at early times. Furthermore, cluster galaxies provide critical tests of the hierarchical paradigm for galaxy formation. In currently popular semi-analytical galaxy formation models in a CDM Universe the descendants of the Ly-break population are massive galaxies in groups and clusters (Baugh et al.\\ 1998). These models also predict that significant differences should exist between early-type galaxies in clusters and those in the general field (Kauffmann 1996). It has been known for a long time that early-type galaxies in clusters form a very homogeneous population: at a given luminosity, they show a very small scatter in their colors, $M/L$ ratios, and line indices (e.g., Bower et al.\\ 1992). The simplest interpretation of this high degree of homogeneity is a small spread in age, although it has been argued that a larger age spread could be ``masked'' by correlated metallicity variations (e.g., Trager et al.\\ 2000)\\footnote{Note that this interpretation requires that we observe early-type galaxies at a special time, when age and metallicity variations exactly cancel.}. Determining the {\\em mean} age of nearby early-type galaxies has proven to be a formidable challenge. The main reason is the well known age-metallicity degeneracy in fitting early-type galaxy spectra (e.g., Worthey 1994). Furthermore, the observed abundance ratios of early-type galaxies cannot be reproduced with simple stellar population synthesis models, which makes absolute determinations of age and metallicity even more uncertain. It is therefore not surprising that most of our understanding of the formation and evolution of cluster early-type galaxies has come from studies of clusters at large lookback times. Since the seminal work by Butcher \\& Oemler (1978) on the colors of galaxies in two distant clusters this field has witnessed great progress: measurements of redshifts, morphologies, colors, $M/L$ ratios, and line indices of cluster galaxies currently span $\\sim 65$\\,\\% of the age of the Universe. Examples of successful ongoing programs are the MORPHS collaboration (Smail et al.\\ 1997), who obtained deep HST images of the cores of $\\sim 10$ clusters at $0.30$. We repeated all our calculations assuming k=1. The results are basically unchanged, with the exception of the temperature interval, which goes from 80 K to 120 K. For comparison, nearby ULIRGs which host an AGN have an effective temperature between 50 and 80 K, when a SED with k$\\sim$1 is assumed. We note however that it is at present extremely difficult to disentangle the starburst and AGN components even in the most studied ULIRGs. It is probable that in many sources hosting an AGN, a starburst contribution is also important. This component can lower the measured effective temperature. The contribution of AGNs to the FIRB strongly affects constraints on galaxy evolution. In background synthesis models in which the stellar component is dominant, for example, a ``passive galaxy evolution'' model cannot reproduce the whole FIRB, unless the stellar initial mass function is strongly different from the Salpeter IMF and is cut at low masses (Franceschini 2000). The significant contribution of at least $\\sim 15$\\% to the FIRB from AGNs indicated by our analysis alleviates this problem, since the total energy emitted by stars and reprocessed in the FIR will be significantly lower by a factor at least 1.7. Alternatively fast galaxy evolution with redshift reproduces the FIR counts in deep surveys with no contribution of AGNs (Chary \\& Elbaz 2001). In this case, as emphasized by these authors, a contribution of AGNs to the FIRB higher than $\\sim$ 15\\%, as we derive, would imply a slower evolution of the galaxy luminosity functions with redshift. Assuming T$_{eff} \\sim 120 K$ we predict a maximum of the AGN contribution from 20\\% to 40\\% at 60~$\\mu$m. If the effective temperature, T$_{eff}$, is higher, the wavelength at which we expect the maximum contribution is lower. In these cases the AGN contribution can easily be higher than 50\\%, and could even reach 100\\%. Again, we remind that if we take into account the absorption at 10-30$~\\mu$m in the most heavily absorbed objects, we predict higher emission at 60-100~$\\mu$m, and a lower one at $\\lambda < 30~\\mu$m. Our argument is not 100\\% watertight: (a) the dust content of the AGNs making the X-ray background could be significantly lower than currently assumed, or (b) the dust composition could be different. \\\\ (a) Several recent studies suggest low dust-to-gas ratio: Granato et al. (1997) propose a low dust-to-gas ratio in X-ray obscured AGNs, from the comparison of the infrared data with their models. Spectroscopically selected quasars suggest the existence of a significant population of AGNs with a low dust-to-gas ratio (Risaliti et al. 2001). As a consequence, the contribution of AGNs to the FIRB estimated in these cases could be negligible: if we assume that the re-radiation band extends down to wavelengths of $\\sim$ 1~$\\mu$m, the FIR emission due to AGNs drop by $\\sim 30$\\%. The drop could be much lower if a significant fraction of the optical-UV primary emission is not reprocessed in the IR. This is for example the case of Broad Absorption Line Quasars, which are heavily absorbed by gas in the X-rays (Mathur, Elvis \\& Singh 1995, Brandt et al. 2000) but little reddened by dust. \\\\ (b) If dust grains in AGNs are much larger, on average, than the standard Galactic composition then the absorption curve of dust would be flatter, and the absorption efficiency significantly lower, than for dust with Galactic size distribution (Laor \\& Draine 1993). A recent study of the continuum and line absorption in nearby intermediate AGNs suggest that this could be the correct scenario for a significant fraction of AGNs (Maiolino et al. 2001a, 2001b). Finally, we note that the obscured AGNs that make the X-ray background may have intrinsically higher X-ray to optical ratio. As a consequence, the bolometric correction we adopted would be overestimated and the total intensity reradiated in the far infrared would be correspondingly lower. This scenario is an alternative to unification models, and cannot be ruled out, given our lack of knowledge on the so-called ``quasar 2s''. In the near future several new observational opportunities are expected, which will tightly constrain our model, revealing the properties and the amount of dust of the AGN contributing to the FIRB: (a) the new X-ray satellites Chandra and XMM-Newton are already providing X-ray spectra of high redshift AGNs. From the comparison between the X-ray and infrared properties of these objects we will be able to have a direct measurement of the contribution of high redshift AGNs to the FIRB. Complementary information, at lower redshift will come from the infrared characterization of the sources found in the deep ASCA and BeppoSAX surveys, like HELLAS (Fiore et al. 1999). (b) Sensitive infrared surveys and identifications, most notably with SIRTF, will test the predictions in Figure 2. More precisely, SIRTF will allow us to measure the SEDs of the hard X-ray background sources, thus giving a direct test for the model proposed here. It will also allow us to study via follow-up of 70~$\\mu$m sources, other contributors to the FIRB (like, for example, Compton thick dusty AGNs). Our model has two main degrees of freedom, as discussed above: the effective temperature of the assumed SED, and the total amount of radiation reprocessed by dust. The optical identification of the sources found in the ISOCAM survey at 15~$\\mu$m will give a first useful constraint. The AGN contribution to the FIRB around $\\sim 60~\\mu$m is quite stable to changes in T$_{eff}$. SIRTF will resolve the bulk of the background at 70~$\\mu$m, reaching a flux of $\\sim 100~\\mu$Jy for point sources, with accurate positions that will enable ready identifications. SIRTF will also resolve a significant fraction of the 160~$\\mu$m (flux limit $\\sim 1$ mJy). These data will tightly constrain our model: a high AGN contribution at 15~$\\mu$m ($>$50\\%) would constrain the SED effective temperature to around 160-180 K. This would imply that the dust responsible to the reprocessing of the optical and UV radiation is much warmer than usually assumed, and that the total amount of dust in AGNs has to be low, with respect to the gas column densities measured in the X-rays. Conversely, a low AGN contribution at 15~$\\mu$m would imply that the dust is on average colder or that the UV and optical primary emission is not almost entirely reprocessed by dust (i.e. there are more dust-free lines of sight than assumed from the X-rays). The 70~$\\mu$m measurement would discriminate between these two scenarios. Finally, if the AGN contribution to the 70~$\\mu$m background is dominant, contrary to our predictions, this would imply that the fraction of completely Compton-thick AGN is much higher than we assumed. Summarizing, in this work we have combined all the available data on AGN SEDs and X-ray background models on one side, and on the FIRB on the other side. We have found, with an approach that is only weakly model-dependent, that the contribution of AGNs to the infrared background is not negligible, and could be dominant in the 15-50~$\\mu$m band (if, for example, the fraction of Compton thick AGNs is higher than our conservative estimate of 20\\%). This has large consequences for models of galaxy evolution. The only possible alternatives are that AGNs have quite different SEDs than usually assumed, or that the dust content or composition in absorbed AGNs is quite different from Galactic. These quantities will be tested soon with ISO and SIRTF." + }, + "0112/astro-ph0112309_arXiv.txt": { + "abstract": "{ Photometric humps in outburst that are locked with the binary orbital period have been observed exclusively in the early phase of outbursts of WZ Sge stars. It is suggested that this \"early hump\" phenomenon is the manifestation of the tidal 2:1 resonance in accretion disks of binary systems with extremely low mass ratios. The \"early humps\" can be understood by the two-armed spiral pattern of tidal dissipation generated by the 2:1 resonance, first discussed by Lin \\& Papaloizou (1979). The tidal removal of angular momentum from the disk during outbursts of dwarf novae, an important feature, is discussed in the context of the disk instability model. The ordering of tidal truncation radius, the 3:1 and 2:1 resonance radius in systems of different mass ratio naturally leads to a classification of dwarf nova systems in three groups according to their mass ratio. The WZ Sge stars are those systems which have the lowest mass ratios and are therefore characterized by \"early humps\". ", + "introduction": " ", + "conclusions": "" + }, + "0112/astro-ph0112145_arXiv.txt": { + "abstract": "We briefly review certain aspects of cosmic microwave background anisotropies as generated in passive and active models of structure formation. We then focus on cosmic strings based models and discuss their status in the light of current high-resolution observations from the BOOMERanG, MAXIMA and DASI collaborations. Upcoming megapixel experiments will have the potential to look for non-Gaussian features in the CMB temperature maps with unprecedented accuracy. We therefore devote the last part of this review to treat the non-Gaussianity of the microwave background and present a method for computation of the bispectrum from simulated string realizations. ", + "introduction": "\\typeout{SET RUN AUTHOR to \\@runauthor} Anisotropies of the Cosmic Microwave Background radiation (CMB) are directly related to the origin of structure in the universe. Galaxies and clusters of galaxies eventually formed by gravitational instability from primordial density fluctuations, and these same fluctuations left their imprint on the CMB. Recent balloon \\cite{boomerang,maxima} and ground-based interferometer \\cite{dasi} experiments have produced reliable estimates of the power spectrum of the CMB temperature anisotropies. While they helped eliminate certain candidate theories for the primary source of cosmic perturbations, the power spectrum data are still compatible with theoretical estimates of a relatively large variety of models, such as $\\Lambda$CDM, quintessence models or some hybrid models including cosmic defects. These models, however, differ in their predictions for the statistical distribution of the anisotropies beyond the power spectrum. The MAP (currently in space) and Planck (scheduled for launch in 2007) satellite missions will provide high-precision data allowing definite estimates of non-Gaussian signals in the CMB. It is therefore important to know precisely what are the predictions of all candidate models for the statistical quantities that will be extracted from the new data, and to identify specific signatures of the various models. There are two main classes of models of structure formation---\\textit{passive} and \\textit{active} models. In passive models, density inhomogeneities are set as initial conditions at some early time, and while they subsequently evolve as described by Einstein-Boltzmann equations, no additional perturbations are seeded. On the other hand, in active models a time-dependent source generates new density perturbations through all time. Most realizations of passive models are based on the idea of inflation. In simplest inflationary models it is assumed that there exists a weakly coupled scalar field $\\phi$, called the inflaton, which ``drives'' the (approximately) exponential expansion of the universe. Quantum fluctuations of $\\phi$ are stretched by the expansion to scales beyond the horizon, thus ``freezing'' their amplitude. Inflation is followed by a period of thermalization, or reheating, during which energy is transferred to usual particles. Because of the spatial variations of $\\phi$ introduced by quantum fluctuations, thermalization occurs at slightly different times in different parts of the universe. Such fluctuations in the thermalization time give rise to density fluctuations. Because of their quantum nature and because of the fact that initial perturbations are assumed to be in the vacuum state and hence well described by a Gaussian distribution, perturbations produced during inflation are expected to follow Gaussian statistics to a high degree, or either be products of Gaussian random variables. This is a fairly general prediction that will be tested shortly with MAP and more thoroughly in the future with Planck data. Active models of structure formation are motivated by cosmic topological defects with the most promising candidates being cosmic strings. It is widely believed that the universe underwent a series of phase transitions as it cooled down due to the expansion. If our ideas about grand unification are correct, then some cosmic defects, such as domain walls, strings, monopoles or textures, should have formed during phase transitions in the early universe \\cite{vilshel}. Once formed, cosmic strings could survive long enough to seed density perturbations. Generically, perturbations produced by active models are not expected to be Gaussian. Defect models possess the attractive feature that they have no parameter freedom, as all the necessary information is in principle contained in the underlying particle physics model. This is a feature that fascinated Dennis Sciama: during his days at Trieste, he would always be willing to discuss these ``marvelous cosmic defects'' with everyone. He even encouraged a few students {\\em sans liaisons contraignantes} to pursue studies on this promising subject and facilitated all necessary means for it. The relevant, ubiquitous dimensionless quantity for every astrophysical and cosmological signature left by strings is $G\\mu$, where $G$ is Newton's constant and $\\mu$ stands for the mass per unit length of the string, with a value of roughly $G\\mu \\sim 10^{-6}$ for GUT strings. This happens to be just the correct order of magnitude of the level of CMB anisotropies. Back in 1996, the power spectrum plot of $l(l+1)C_l$ vs.~$l$ consisted of some scattered points and strings were considered a likely candidate for the origin of structure formation. Moreover, the error bars were so big that Dennis would never miss the opportunity to tease inflation fans who wanted to see a peak in that plot: ``Come on'', he would say, ``you do not claim you actually {\\em see\\/} a peak there, do you..!''. To close this section, let us mention that, in addition to purely active or passive scenarios, perturbations could be seeded by some combination of the two mechanisms. For example, cosmic strings could have formed just before the end of inflation and partially contributed to seeding density fluctuations. It has been shown \\cite{hybrid} that such hybrid models are very successful in fitting the CMB power spectrum data. Cosmic strings are generally expected to produce distinguishing non-Gaussian features in the CMB and it will soon become possible to look for them in the data from MAP and Planck. ", + "conclusions": "" + }, + "0112/astro-ph0112373_arXiv.txt": { + "abstract": "We examined the VSNET light curve of the ER UMa-type star V1159 Ori. We detected a large variation of the supercycle (the interval between successive superoutbursts) between extremes of 44.6 and 53.3 d. The outburst activity was also found to decrease when the supercycle was long. The observed variation of the supercycle corresponds to a variation of $\\sim$40\\% of the mass-transfer rate from the secondary star, totally unexpected for this class of objects. We also detected a hint of $\\sim$1800 d periodicity in the variation, whose period is close to what has been suggested for solar-type cycles for cataclysmic variables (CVs). If this periodicity is caused by the magnetic activity of the secondary star, this detection constitutes the first clear evidence of continuing magnetic activity in CV evolution, even after crossing the period gap. This activity may partly explain still poorly understood origins of the high mass-transfer rates in ER UMa-type stars. ", + "introduction": "ER UMa stars are a subgroup of SU UMa-type dwarf novae (for a review of dwarf novae, see \\cite{osa96}), whose known members are ER UMa, V1159 Ori, RZ LMi, and DI UMa. The most striking feature of ER UMa stars is the extremely short recurrence time (19--45 d) of superoutbursts (\\cite{kat95}; \\authorcite{nog95a} 1995a,b; \\cite{rob95}; \\cite{mis95}; \\cite{kat96}). Another striking feature of ER UMa stars is the stability of supercycles, both in their lengths and outburst pattern. The best exemplification of this stability can be seen in folded light curves and $O-C$ figures presented in \\citet{rob95}. The extremely short supercycle length and the stability of the outburst patterns are basically explained, within the framework of the disk-instability model, as a result of constant high mass-transfer rates from the secondary \\authorcite{osa95a} (\\yearcite{osa95a}). The mass-transfer rates in SU UMa-type dwarf novae are generally considered to be confined to a small range determined by angular-momentum removal by the gravitational wave radiation. The origin of high-mass transfer rates in ER UMa stars is still an open question. Some models assume irradiation effect from a hot white dwarf, which may be the result of a hypothetical recent nova eruption (the possibility was originally raised by \\authorcite{nog95b} \\yearcite{nog95b}, see also \\cite{pat98}). An examination of any secular changes in the supercycle in these systems would provide an essential clue for testing these hypotheses. ", + "conclusions": "The long-term average of the supercycle lengths in V1159 Ori, being close to the minimum value predicted by \\authorcite{osa95a} (\\yearcite{osa95a}), the supercycle length near this period is expected to be insensitive to the mass-transfer rate from the secondary. If the observed change in V1159 Ori was caused by the variable mass-transfer rate, a relatively large change is necessary to reproduce the observation. Using the $\\dot{M}-{\\rm supercycle}$ diagram in \\authorcite{osa95a} (\\yearcite{osa95a}), a supercycle of 53.3 d corresponds to a reduction of $\\sim$40\\% of mass-transfer rates from what is expected for a 44.5-d supercycle. The marked reduction of the superoutburst duty cycle during this period (figure \\ref{fig:figure3}) also supports this interpretation. Another observational evidence of a large period change in ER UMa stars has been reported in DI UMa \\citep{fri99}. However, the extreme shortness of supercycles in DI UMa and RZ LMi requires an additional (still poorly identified) mechanism (\\authorcite{osa95a} \\yearcite{osa95a}), and its change may be of different nature. Another noteworthy feature in the observed $O-C$ diagram of V1159 Ori is a possible periodicity with a period of $\\sim$38 cycles, corresponding to $\\sim$1800 d, rather than a monotonous change originally proposed by \\citet{rob95}; this is contrary to the expected effect by decreasing heating from a hypothetical recent nova eruption on a white dwarf. The observed possible long-term period is close to those observed as possible solar-type cycles in cataclysmic variables (e.g. \\cite{bia88}; \\cite{ak01}). If such a ``solar-type\" cycle is responsible for the change in the supercycle of V1159 Ori, this may provide promising evidence for the presence of magnetic activity in dwarf novae below the period gap, which has usually been considered to cease, or to be markedly reduced, when the secondary becomes fully convective after crossing the period gap. Furthermore, the continuing magnetic activity may be one of the mechanisms for effectively removing the angular momentum from the binary system, by which the required high mass-transfer in ER UMa-type systems may be partly explained. \\vskip 3mm The author is grateful to VSNET members, especially to Rod Stubbings, Gene Hanson, Gary Poyner, Andrew Pearce, Seiichiro Kiyota, Eddy Muyllaert, Tsutomu Watanabe and numerous observes for providing vital observations." + }, + "0112/astro-ph0112529_arXiv.txt": { + "abstract": "We present and analyze the optical and X-ray catalogs of moderate-redshift cluster candidates from the ROSAT Optical X-ray Survey, or ROXS. The survey covers the sky area contained in the fields of view of 23 deep archival ROSAT PSPC pointings, 4.8 square degrees. The cross-correlated cluster catalogs were constructed by comparing two independent catalogs extracted from the optical and X-ray bandpasses, using a matched-filter technique for the optical data and a wavelet technique for the X-ray data. We cross-identified cluster candidates in each catalog. As reported in Paper I, the matched-filter technique found optical counterparts for at least 60\\% (26 out of 43) of the X-ray cluster candidates; the estimated redshifts from the matched filter algorithm agree with at least 7 of 11 spectroscopic confirmations ($\\Delta z \\lesssim 0.10$). The matched filter technique, with an imaging sensitivity of $m_I \\sim 23$, identified approximately 3 times the number of candidates (155 candidates, 142 with a detection confidence $>3\\sigma$) found in the X-ray survey of nearly the same area. There are 57 X-ray candidates, 43 of which are unobscured by scattered light or bright stars in the optical images. Twenty-six of these have fairly secure optical counterparts. We find that the matched filter algorithm, when applied to images with galaxy flux sensitivies of $m_I \\sim 23$, is fairly well-matched to discovering $z\\leq1$ clusters detected by wavelets in ROSAT PSPC exposures of 8,000-60,000 seconds. The difference in the spurious fractions between the optical and X-ray (30\\% and 10\\% respectively) can not account for the difference in source number. In Paper I, we compared the optical and X-ray cluster luminosity functions and we found that the luminosity functions are consistent if the relationship between X-ray and optical luminosities is steep ($L_x \\propto L_{opt}^{3-4}$). Here, in Paper II, we present the cluster catalogs and a numerical simulation of the ROXS. We also present color-magnitude plots for several of the cluster candidates, and examine the prominence of the red sequence in each. We find that the X-ray clusters in our survey do not all have a prominent red sequence. We conclude that while the red sequence may be a distinct feature in the color magnitude plots for virialized massive clusters, it may be less distinct in lower-mass clusters of galaxies at even moderate redshifts. Multiple, complementary methods of selecting and defining clusters may be essential, particularly at high redshift where all methods start to run into completeness limits, incomplete understanding of physical evolution, and projection effects. ", + "introduction": "Clusters of galaxies are the most massive gravitationally-bound systems in the universe. Because they sample the high mass end of the mass function of collapsed systems, they can be used to determine cosmological parameters such as $\\Omega_{m}$ (e.g. Donahue \\& Voit 1999). The largest clusters ($\\sim 10^{15} ~M_\\odot$), are the products of the collapse of matter from a very large volume of space ($r \\sim 16 h^{-1}(\\Omega/0.2)^{-1/3} ~\\rm{Mpc}$). Therefore they are thought to be ``fair samples'' of the universe -- that the mass to light ratio or the baryonic mass fraction defined within the domain of a cluster of galaxies is representative of that ratio in the universe as a whole. They are purported to be ``closed boxes'' to star formation and evolutionary processes that occur within their domain. In this paper we present and analyze the catalogs from our joint optical-X-ray search for clusters of galaxies. We conducted this survey in order to provide a sample of clusters to test such assumptions about clusters of galaxies and to investigate the impact of sample selection on studies of cluster evolution and the evolution of their member galaxies. Understanding cosmological or galaxy evolution studies of clusters critically requires an understanding of the biases in any sample of clusters. For example, the evolution of the number density of systems with cluster-sized masses as a function of mass and redshift is a fundamental prediction of cosmological structure formation models. To know the number density of clusters, we must know the biases inherent in how we find them, preferably as a function of cluster mass. Furthermore, testing the ``fair sample'' hypothesis requires reliable and unbiased selection of the most massive clusters. Ever since Abell (1958) and Zwicky (Herzog, Wild \\& Zwicky 1957; Zwicky 1961) began publishing catalogues of optically selected clusters of galaxies, the definition of a cluster and the definition of biases inherent in the cluster detection process have been lively topics of debate. In 1978, the launch of the first X-ray imaging telescope, the Einstein observatory, began a new era of cluster discovery, as clusters proved to be luminous ($>10^{42-45} \\lum$), extended ($r\\sim1-5$ Mpc) X-ray sources, readily identified in the X-ray sky. The intracluster gas, in nearly hydrostatic equilibrium with the gravitational potential of the cluster, radiates optically thin thermal bremstrahlung and line radiation. X-ray selection of clusters is more robust against contamination along the line of sight than traditional optical methods since the richest clusters are relatively rare and since X-ray emissivity, which is proportional to the gas density squared, is far more sensitive to physical overdensities than is the projected number density of galaxies on the sky. One cluster sample differs from another depending on how the clusters were detected. Optical selection of clusters using traditional methods looking for overdensities of galaxy counts (e.g. Abell 1958) was rife with contamination problems. However, modern methods such as the ``matched filter'' algorithm (Postman et al.\\ 1996, P96 hereafter) provide automated, uniform detection of galaxy overdensities in deep optical images. The matched filter technique searches for local density enhancements in which galaxies follow a magnitude distribution characteristic of that expected for a cluster of galaxies. The results include statistically quantifiable estimates of cluster richness, redshift, and significance. The first X-ray selection methods using sliding boxes were optimal for point sources. Thus, the detection method used to construct the Extended Medium Sensitivity Survey (EMSS; Gioia et al. 1990b) was biased somewhat towards selecting clusters with high central surface brightnesses. Now there are several algorithms optimized for detecting extended sources, including wavelets (Rosati et al. 1995) and Voronoi-Tesselation Percolation methods (Scharf et al. 1997). A decade ago, optical and X-ray surveys apparently disagreed about how much clusters have evolved since $z\\sim0.5-1.0$. Optical surveys indicated very little evolution since $z\\sim0.5-1$ (Gunn, Hoessel \\& Oke 1986), but the accurate measurements of survey volumes and cluster properties required for quantitative assessment of this evolution were difficult to quantify in these first high-redshift cluster surveys and the volumes were small so uncertainties were large. X-ray studies suggested modest evolution (Gioia et al. 1990a). The most recently compiled X-ray samples of clusters over a range of redshifts out to $z\\sim0.8-1.2$ agree that the X-ray luminosity function for moderate luminosity clusters has not evolved significantly since $z\\sim0.8$ (Borgani et al. 1999; Nichol et al. 1999; Rosati et al. 1998, 2000; Jones et al. 1998), while the most luminous systems, contained in the EMSS, might have evolved somewhat (Henry et al. 1992; Nichol et al. 1997, Vikhlinin et al 1998, 2000; Gioia et al. 2001) or very little (Lewis et al. 2002). More recent optical surveys for distant clusters continue to find very little evidence for cluster number density evolution at moderate redshifts (Couch et al. 1991; P96). The explanation of what may seem like a persisting discrepancy is that if any evolution exists in the X-ray cluster population, it is only occuring in the highest luminosity systems which are also the most rare systems. The optical surveys of Couch et al. (1991) and P96 were too small and shallow to detect the putative evolution of the rarest systems. While the most recent optical and X-ray results are now at least in statistical agreement on the question of evolution since $z<0.8$, the question remains whether both techniques are selecting the same clusters. The fundamental quantity, from the viewpoint of comparison to cosmological simulations, is the cluster's mass. We do not know {\\em a priori} whether optical luminosity or X-ray luminosity should be better correlated with a cluster's mass. The fundamental question, from the viewpoint of ``fair sample'' techniques of measuring universal ratios, is whether clusters are truly a ``fair sample''. For example, M/L ratios depend on the bandpass of the light and the star formation history of the constituent galaxies. If the gas fractions or the M/L ratios vary significantly from cluster to cluster they are obviously not representative of the universe as a whole. X-ray selection is generally thought to be superior to optical selection. Observationally, the hot gas is a larger fraction of the cluster mass than the stellar mass, and the X-ray luminosity of a cluster is far easier to measure than its optical luminosity. For X-ray selected clusters, studies of gas fractions and cluster M/L ratios show that these quantitites are statistically constant (Evrard 1997; Arnaud \\& Evrard 1999; Carlberg et al. 1996). However, if X-ray selection biases the selection of the clusters, high-mass clusters of galaxies with low hot gas fractions (if they exist) would be omitted from such studies. Massive clusters, under the ``fair sample'' hypothesis, should have nearly identical baryon fractions, however they are discovered. With the ROSAT Optical X-ray Survey (ROXS) for clusters of galaxies, we have endeavored to address such issues by obtaining optical images of complete 30' by 30' fields centered on positions of deep ROSAT PSPC pointings. In contrast to previous ROSAT PSPC serendipitous surveys such as those conducted by Rosati et al. (1995, 1998), Jones et al. (1998), Romer et al. (2000), and Vikhlinin et al. (1998), the ROXS includes optical imaging for the entire field of view of each X-ray pointing. The X-ray selection and optical selection of cluster candidates was then done independently of each other. We observed 23 ROSAT pointings for a total of nearly 5 square degrees in I band. For five of these fields we also obtained V-band imaging. In this paper (Paper II) we present the catalogs, survey windowing functions, data reduction and observation details, an analysis of detection likelihoods, as well as an expanded discussion and further analysis, including numerical simulations of the survey. In \\S2, we describe the X-ray field selection criteria and the optical observations. In \\S3, we present the optical cluster candidates catalog, cross-identification of clusters in the V and I bands. In \\S4, we present the X-ray cluster candidate catalogs and the X-ray/optical cross-identification procedure. In \\S5 we describe properties of the cluster candidates, including the distribution of observed properties of objects in the sample, the estimated richness vs. $L_x$, the $V-I$ vs $I$ color magnitude diagrams for the clusters identified in both the $V$ and $I$ bands. We discuss and summarize our results in \\S6 and \\S7 respectively. For all derived quantities, we have used $H_0=75 h_{75} $ km s$^{-1}$ Mpc$^{-1}$, and $q_0=0.5$. ", + "conclusions": "The optical matched-filter selection technique works well to find candidate clusters of galaxies, but the spurious fraction is high at $\\sim30\\%$, and it is demonstrably not 100\\% complete - it misses clusters at a rate of 10-20\\% at least. The selection window for a matched filter cluster sample is more difficult to quantify because of its sensitivity to the spectral energy distribution of the galaxies for which it is searching. On the other hand, the X-ray selection technique, while generating fewer spurious cluster candidates ($\\sim10\\%$) and reliably finding the (apparently) more massive clusters, misses some high $\\Lambda_{cl}$ cluster candidates. We do not know yet whether these are true massive clusters or whether they are fortuitous projections of less massive systems. Redshifts and X-ray observations are required to determine the nature of these candidates. Zabludoff \\& Mulchaey (1998) and Mulchaey \\& Zabludoff (1998) showed in their sample of 12 optically selected groups that groups without detected X-ray emission tend to be lower velocity dispersion systems with few or no elliptical galaxies. The undetected groups in the Zabludoff \\& Mulchaey sample may not even be bound. Some X-ray observations of optically-selected, high redshift clusters have revealed such clusters to be ``underluminous'' in the X-ray (Castander et al 1994.) The high-significance, optically rich ROXS candidates without X-ray counterparts are prime targets for observational tests whether X-ray selection indeed selects on the basis of cluster mass. The correlation of the presence of a dense, X-ray emitting intracluster media and the presence of a significant elliptical galaxy population is another testable hypothesis with further observations of the ROXS sample. With the limited color information we have in hand, we were not able to show that X-ray cluster candidates at the moderate X-ray luminosity levels available in our sample are more likely to have red elliptical sequences. But we were only able to make this test for a small number of cluster candidates, given the available photometry. A related result from ROXS is that the cluster richness $N_R$ is not well-correlated with cluster X-ray luminosity, a result which is supplementary to our result in Paper I regarding cluster X-ray luminosity and the $\\Lambda_{cl}$ richness parameter. The clusters in our sample are on average poorer and less massive than those found in a half-sky survey (e.g. XBACS), owing to the smaller sky area and greater depth of our survey. For $\\Lambda > 30$, Donahue et al. (2001) suggested that the joint redshift and $\\Lambda_{cl}$ distribution arises from the population of X-ray clusters with a steep dependence between $L_x$ and $\\Lambda_{cl}$. Equivalently, we show here that we can also reproduce this joint distribution with a large scatter between $L_x$ and $\\Lambda_{cl}$ in Monte-Carlo simulations of the ROXS. A correlation between cluster X-ray luminosity and cluster richness has been suspected at least since Jones \\& Forman (1978) plotted cluster richness class against X-ray luminosity for nearby ($z<0.07$) clusters. But even in their sample, a $3 \\times 10^{44} ~\\lum$ cluster is equally as likely to have a cluster richness of 0 as it is 2 or 3. Only the most luminous clusters ($L_x > 10^{45} ~\\lum$) had cluster richnesses reliably of 2 or 3. The XBACS clusters (Figure~\\ref{LXrich}) show this correlation and very large scatter with cluster richness. However, richness has long been a suspect observational parameter for X-ray astronomers (e.g. Mushotzky et al 1978), since the measurements rely on number counts that only become more difficult to ascertain at higher redshift with the associated higher contamination levels. The matched filter method provides a possibly more robust and objective means of estimating a cluster candidate's luminosity in the form of $\\Lambda_{cl}$, yet, this measure too is not strongly correlated with cluster X-ray luminosity. The lack of strong correlation between X-ray and optical luminosity, and the large scatter of those two quantities, are confirmed by the ROX Survey. A review by Borgani \\& Guzzo (2001) showed that the velocity dispersions and cluster optical luminosities of an optically-selected sample of clusters are not as well-correlated as the velocity dispersions and the optical luminosities of a subsample of clusters selected for their X-ray fluxes from the original sample. The velocity dispersions of these X-ray selected clusters are even better correlated with their X-ray luminosities, suggesting that X-ray luminosity is better correlated with cluster mass than is optical luminosity. The ROX Survey has no such independent test of correlation with mass; however, followup observations of this sample in the X-ray and the optical will further test the Borgani \\& Guzzo suggestion at higher redshifts. The ROXS, since it is a uniformly selected sample at both X-ray and optical wavelengths, has examples of both X-ray candidates without optical counterparts and optical candidates without X-ray counterparts. Both subsamples probe the extremes of the $L_x/L_{opt}$ distribution. Explaining the extremes will go a long way towards explaining why and how the emissivity of the intracluster gas is related to the light emitted by the stars in the cluster. It is particularly important to see if there is a ``third parameter\" such as galaxy formation efficiency, evolutionary stage of the cluster, or galaxy morphology populations, that creates the large spread in the $L_x$ versus richness relationship for clusters, or if the large spread is due to projection effects significantly impacting the optical selection technique. One possible clue is the existence of a difference in the conclusions based on cluster contents from an optically-selected, rather heterogenous sample of clusters (the MORPHS sample, in Smail et al 1997) and those found in an X-ray selected sample of clusters by Ellingson et al (2001). The X-ray selected clusters show cores that exhibit very little evolution between $z=0.5$ and the present, whereas the MORPHS clusters show evidence that the S0 population may be turning into ellipticals during that same time. It is possible that the presence of a well-developed intracluster medium is ubiquitous in a massive cluster, but in less massive clusters, the galaxy populations are still evolving in the cluster cores. Uniform morphological studies of a cluster sample diverse in X-ray and optical properties are needed to test such a statement, and such a test could be accomplished with a sample such as the ROXS. Rosati et al. (1995, 1998) selected clusters for the RDCS not only for their X-ray emission, but used an extent criteria as well. There was some worry that this selection criteria may have missed the most compact X-ray clusters which remain unresolved by the ROSAT PSPC, particularly those at high redshift. However, our survey does not reveal an obvious population of high redshift cluster candidates with X-ray point sources. The correspondence between an optical cluster and an X-ray point source is not more than we would expect by chance - therefore we do not believe we have found a population of cluster sources which would go undetected by the ROSAT cluster surveys which select for extent, a result consistent with the findings of the Wide Angle ROSAT Pointed Survey (WARPS; Jones et al 1998), a ROSAT serendipitous search which did not filter for extent. We have some examples of optical clusters with X-ray point source counterparts which could be AGN in those clusters. Approximately 25 of our candidates have X-ray point sources within 1'-2' of the cluster centroid; optical identification of these sources is required to ascertain whether the candidate and the point source are physically related. Probably the most surprising result is that the ROXS cluster candidates do not, as a rule, show distinct or strong red sequences in the color-magnitude diagrams of their galaxy populations. Only about 50\\% of the optical cluster candidates detected in both I- and V-band show galaxy colors distinct from those of the field population. And only $\\sim50\\%$ of the X-ray clusters so identified show such sequences. It is possible that the red sequence is ubiquitous in the most massive, virialized clusters. But perhaps at some unknown threshold in mass or level of virialization or luminosity, the red sequence ceases to be distinct and prominent. We plan to use the red sequence detection method (Gladders \\& Yee 2000), which does not rely on an obvious red sequence to select cluster candidates, to see what cluster candidates it finds in the ROXS galaxy photometry data. The ROXS sample identifies several problems that warrant further investigation: \\begin{itemize} \\item A red sequence is not prominent in all X-ray selected clusters. At what X-ray luminosity (or equivalently at what mass) does the red sequence cease to be prominent? Is the presence of a red sequence dependent on mass? \\item The estimated optical luminosity and richness of a cluster are not at all well-correlated with the X-ray luminosity. There are at least three unsolved issues regarding the correlation of the optical luminosity, X-ray luminosity, and cluster virial mass: \\begin{enumerate} \\item Is the lack of correlation representative of the difficulties of estimating the optical luminosity or is it representative of an intrinsic variation or scatter in the M/L ratios in clusters? \\item If the estimated optical or X-ray luminosity, and therefore the detectability of a cluster, has enormous scatter with respect to the mass of the cluster, selection by mass using optical or X-ray light may be very difficult. Knowing the scatter with respect to cluster mass with both optical and X-ray luminosity is important. Individual studies show good correspondence of both quantities with velocity dispersion (e.g. Girardi et al. 2000 and Mahdavi \\& Geller 2001 for $L_B$ and $L_x$ respectively), but a comparative analysis of the masses and luminosities of a well-chosen, homogeneous sample has not been done yet. \\item If there is such a large intrinsic variation in the optical M/L for clusters, what evolutionary process controls this variation? Can galaxy formation efficiency or evolutionary history be so radically different from one cluster environment to another to produce this effect? Does M/L correlate with any other property of a cluster of galaxies? \\end{enumerate} \\item While galaxies and hot gas may trace the gravitational potentials of the most massive clusters, mergers and other physics may disrupt that correspondence. True optical condensations of galaxies with low levels of X-ray emission may be galaxies which have been temporarily separated from the hot intracluster gas while undergoing a merger event; galaxies experience the merger as a non-collisional fluid while the gas experiences hydrodynamic effects of shocks and cooling. An observational example of such a merger is Abell~754 (Zabludoff \\& Zaritsky 1995). Studies of the ROXS ``X-ray poor'' cluster candidates may reveal their true nature, whether they are mergers, minor systems embedded in filaments, or mere projection effects. \\end{itemize} Because the ROXS has greater depth and smaller sky area than surveys like the EMSS or XBACS, the ROXS clusters must be relatively low-mass clusters compared to typical clusters in those surveys. The low-mass end of the cluster mass function may be where the physics of the intracluster medium collides with the physics of star formation and galactic winds. The energy injected by stellar processes is closer to the specific energy per particle in the gravitationally bound gas of a low mass cluster of galaxies. Mergers may be more common or at least more significant in the low-mass clusters -- when the relative velocities of the galaxies are not much greater than the internal velocities of the galaxies, collisions and interactions are more likely to induce star formation. The least massive clusters may be the most recent members to the cluster hierarchy, whose galaxies are most recently accreted from the field. The optical and X-ray properties of such clusters may thus still be correlated, but only very weakly and with high scatter. Such recent formation may explain the weakness or lack of a red sequence, which may only dominate in clusters where the oldest ellipticals formed preferentially in high density regions at an early time. The other inference possible from ROXS is that at some point below a given mass on the mass scale of clusters, some treasured assumptions about clusters, their M/L constancy, the invariance of the baryon fraction in clusters, and their ubiquitous red galaxy content may break down." + }, + "0112/astro-ph0112003_arXiv.txt": { + "abstract": "The young and energetic pulsar \\psr\\ powers a bright X-ray synchrotron nebula, embedded in the unusual supernova remnant \\snr. We present observations of this system with the {\\em Chandra X-ray Observatory}, which show a spectacularly complicated source. The nebula is dominated by a bright collimated feature which we interpret as a relativistic jet directed along the pulsar spin axis. Several compact knots can be seen in the immediate vicinity of the pulsar. While many of these features are similar to those seen around the Crab pulsar, the nebula surrounding PSR~\\psr\\ shows important differences which are possibly a result of the the latter's low nebular magnetic field and low density environment. ", + "introduction": "The supernova remnant (SNR) \\snr\\ (MSH~15--5{\\em 2}) has unusual radio and X-ray properties. At radio wavelengths (Figure~\\ref{fig_1}), its morphology is dominated by two distinct limbs of radio emission. Superimposed on the northern limb is a bright core of emission, coincident with the H$\\alpha$ nebula RCW~89. The distinctly non-circular appearance of this source has been attributed to expansion into an elongated cavity, a claim supported by recent \\HI\\ observations of this region (Dubner \\etal\\ 2002). \\HI\\ absorption towards the SNR demonstrates it to be at a distance of 5~kpc (Gaensler \\etal\\ 1999). At X-ray energies (Figure~\\ref{fig_1}), the system is dominated by a bright central point source. This corresponds to the pulsar \\psr, which has also been detected at radio wavelengths and in $\\gamma$-rays. PSR~\\psr\\ is one of the youngest and most energetic pulsars known: it has a spin-period $P=151$~ms, a magnetic field $B = 1.5\\times10^{13}$~G, a spin-down luminosity $\\dot{E} = 1.8\\times10^{37}$~erg~s$^{-1}$ and a characteristic age $\\tau = 1700$~yr. Surrounding the pulsar is an elongated non-thermal nebula, presumed to be the pulsar wind nebula (PWN) powered by the pulsar's spin-down; no radio counterpart to this PWN has been identified. To the north of PSR~\\psr\\ is a source of thermal X-rays, coincident with the bright radio and optical emission from RCW~89 (Trussoni \\etal\\ 1996; Tamura \\etal\\ 1996). \\begin{figure} \\centerline{\\psfig{file=gaenslerb1_1.eps,height=8cm,angle=270}} \\caption{A radio/X-ray comparison of \\snr. The greyscale corresponds to 843-MHz MOST observations, while the white contours represent smoothed {\\em ROSAT}\\ PSPC data. The position of PSR~\\psr\\ is marked with a ``+'', while the black box delineates the ACIS-I field-of-view.} \\label{fig_1} \\end{figure} Existing observations have raised a number of issues regarding this pulsar and its interaction with its environment. Firstly, the elongated morphology of the PWN suggests that its morphology is dominated by a collimated outflow directed along the pulsar spin-axis, which possibly collides with and is interacting with the RCW~89 region (Manchester \\& Durdin 1983; Tamura \\etal\\ 1996; Brazier \\& Becker 1997; Gaensler \\etal\\ 1999). Greiveldinger \\etal\\ (1995) have claimed that there is a compact disc of nebular emission immediately surrounding the pulsar, while Brazier \\& Becker (1997) rather propose a ``cross''-shaped morphology in this region, which they interpret as an equatorial torus and polar jets, seen edge-on. Clearly our understanding of this complicated source can benefit from observations at higher angular resolution. We have consequently carried out observations of PSR~\\psr\\ and its surroundings with the {\\em Chandra X-ray Observatory}. We summarize these results below; these data are discussed in more detail by Gaensler \\etal\\ (2002). ", + "conclusions": "These \\cxo\\ data have provided a wealth of new information on PSR~\\psr\\ and its interaction with its environment. We have confirmed the presence of a collimated flow directed along the pulsar spin-axis, and have argued that the flow is relativistic and is inclined at $<30^\\circ$ to our line-of-sight. We have interpreted two arcs of emission seen close to the pulsar as dynamical features in an equatorial flow, and have identified several knots at separations $<0.5$~pc from the pulsar. Many issues still need to be investigated. Do the arcs and knots show motion and/or variability? Do any of these features have counterparts at other wavelengths? What is the nature of RCW~89 and its thermal clumps? While these questions still remain, it is clear that PSR~\\psr\\ provides a new opportunity to probe the detailed structure of a pulsar wind." + }, + "0112/gr-qc0112044_arXiv.txt": { + "abstract": "We consider simple hydrodynamical models of galactic dark matter in which the galactic halo is a self-gravitating and self-interacting gas that dominates the dynamics of the galaxy. Modeling this halo as a sphericaly symmetric and static perfect fluid satisfying the field equations of General Relativity, visible barionic matter can be treated as ``test particles'' in the geometry of this field. We show that the assumption of an empirical ``universal rotation curve'' that fits a wide variety of galaxies is compatible, under suitable approximations, with state variables characteristic of a non-relativistic Maxwell-Boltzmann gas that becomes an isothermal sphere in the Newtonian limit. Consistency criteria lead to a minimal bound for particle masses in the range $30 \\,\\hbox{eV} \\leq m \\leq 60 \\,\\hbox{eV}$ and to a constraint between the central temperature and the particles mass. The allowed mass range includes popular supersymmetric particle candidates, such as the neutralino, axino and gravitino, as well as lighter particles ($m\\approx$ keV) proposed by numerical N-body simulations associated with self-interactive CDM and WDM structure formation theories. ", + "introduction": "\\noindent The presence of large amounts of dark matter at the galactic lengthscale is already an established fact. Assuming that this dark matter is a gas (or gas mixture) of various particles species, the established classification criteria labels possible dark matter forms as ``cold'' or ``hot'', depending on the relativistic or non-relativistic nature of the particles' energetic spectrum at their decoupling from the cosmic mixture \\cite{Trimble}, \\cite{tbook_1}, \\cite{tbook_2}. Hot dark matter (HDM) scenarios seem to be incompatible with current theories of structure formation and thus, are not favoured dark matter candidates \\cite{tbook_1}, \\cite{tbook_2}, \\cite{dm_scenarios}. Cold dark matter (CDM), usualy examined within a newtonian framework, can be considered as non-interactive (a self gravitating gas of collisionless particles) or self-interactive \\cite{old_cdm_theo}. CDM models are often developed in terms of n-body numerical simmulations \\cite{old_cdm_nbody}, \\cite{nbody_1}, \\cite{nbody_2}, \\cite{nbody_3}. Non-interactive CDM models present the following discrepancies with observations at the galactic scale \\cite{cdm_problems_1}, \\cite{cdm_problems_2}: (a) the ``substructure problem'' related to excess clustering on sub-galactic scales, (b) the ``cusp problem'' characterized by a monotonic increase of density towards the center of halos, leading to excessively concentrated cores. These problems appear in the more recent numerical simulations (see \\cite{nbody_1}, \\cite{nbody_2}, \\cite{nbody_3}). In order to deal with these problems, the possibility of self-interacting dark matter has been considered, so that nonzero pressure or thermal effects can emerge, thus leading to self-interactive models of CDM ({\\it i.e.} SCDM) \\cite{scdm_1}, \\cite{scdm_2}, \\cite{scdm_3}, \\cite{scdm_4}, \\cite{scdm_5} and ``warm'' dark matter (WDM) models \\cite{wdm_1}-\\cite{wdm_6} that challenges the duality CDM vs. HDM. Other proposed dark matter sources consist replacing the gas of particles approach by scalar fields \\cite{sfe_1}, \\cite{sfe_2} and even more ``exotic'' sources \\cite{sfe_3}. Whether based on SCDM or WDM, current theories of structure formation point towards dark matter characterized by particles having a mass of the order of at least keV's (see \\cite{scdm_1}-\\cite{wdm_6}), thus suggesting that massive but light particles, such as electron neutrinos and axions (see Table~1), should be eliminated as primary dark matter candidates (though there is no reason to assume that these particles would be absent in galactic halos). Of all possible particle candidates (denoted as WIMP's: weakly interactive massive particles) complying with the required mass value of relique gases, only the massive neutrinos (the muon or tau neutrinos), have been detected, whereas other WIMPS (gravitino, sterile neutrino, axino, etc.) are speculative. See \\cite{Pal}, \\cite{SNO}, \\cite{PDG}, \\cite{Ellis} and Table~1 for a list of candidate particles and appropriate references. Even if the dynamics of visible matter in galaxies can be described succesfuly with Newtonian gravity, we believe that General Relativity is an appropriate framework for understanding basic features of galactic dark matter, a gravitational field source whose precise physical nature still remains an open question. If the results obtained with GR coincide with Newtonian results, then there is no harm done from a pragmatic calculations-oriented point of view. However, from a formal-theoretical approach, we believe it is beneficial to broaden the scope of the study of galactic dynamics by incorporating it to a more general gravitational theory. In particular, in this paper we aim at testing the compatibility between observed galactic rotation curves and simple thermodynamical assumptions under the framework of GR. Since dark matter halo probably constitutes the overwhelming mayority (90 \\%) of the galactic mass, an alternative approach to numerical simulations and newtonian hydrodynamics follows by a general relativistic model describing the gravitational field of the galaxy as a spacetime geometry generated by the the dark matter halo (as a self gravitating gas), hence visible matter becomes test particles that evolve along stable geodesics of this spacetime. There is strong empiric evidence that the radial profile of rotational velocities (``rotation curves'') in most galaxies roughly fits a ``universal rotation curve'' (URC) \\cite{urc_1}, \\cite{urc_2}. This URC is characterized by a ``flattening'' effect whereby rotation velocities tend to a constant ``terminal'' velocity whose value depends on the type of galaxy (between $125\\,\\hbox{km/sec}$ and $250\\,\\hbox{km/sec}$). The profile of rotation curves identifies two main contributions of galactic matter: visible matter (the disk), showing a keplerian decay, and dark matter (the halo), explaining the flattening effect. This kinematic evidence might allow us to determine (at least partialy) the geometry of the spacetime associated with a self gravitating galaxy. In other words, our approach somehow inverts the standard initial value procedure in general relativistic hydrodynamics: instead of prescribing initial data based on physicaly motivated sources and then find the geometry of spacetime and the trayectories of test particles after solving Einstein's equations, we provide first constraints on the geometry of spacetime (from symmetry criteria and empirical kinematic data) and then find, with the help of the field equations, the corresponding momentum-energy tensor of the sources. This approach to galactic dark matter has been used in connection to scalar fields \\cite{sfe_1}. Bearing in mind that the dark matter halo overwhelmingly dominates the galactic matter content (at least in the halo region), we shall assume that the galactic halo (as a self gravitating gas) is the unique matter source of the galactic spacetime. Visible matter becomes then test observers that follow stable circular geodesic orbits (the galactic rotation curves) of this spacetime. Following the ``inverse'' approach described above, we propose to use the empiric law governing the form of the URC for the galactic halo (see \\cite{urc_1} and \\cite{urc_2}) in order to make specific asertions on the nature of the sources of the galactic spacetime. Considering the self-gravitating galactic halo gas to be self-interactive (instead of colissionless matter or a scalar field), we aim at verifying if the assumption of the URC profile for the rotation velocity of geodesic observers (rotation curves) is compatible with the assumption that the galactic halo gas is a simple self-gravitating and self-interactive gas in thermodynamical equilibrium. For this purpose, we consider the galactic halo to be a spacetime characterized as: (a) sphericaly symmetric, (b) its energy-momentum tensor is that of a perfect fluid satisfying the equation of state of an equilibrium Maxwell-Boltzmann gas in its non-relativistic limit \\cite{RKT}, \\cite{rund}. Assumption (a) is supported by observations in the halo of galaxies, while (b) is the central hypothesis in the present work. Assumption (b) requires spacetime to be stationary (static if rotation vanishes) and leads to a law relating temperature gradients with the 4-acceleration (Tolman Law). Then, since we are assuming the validity of the empiric URC for the galactic halo, we need to cast the field equations and the conditions imposed by the thermodynamics in terms of this rotation velocity ({\\it i.e.} the velocity of test particles in stable geodesic orbits), considered now as a dynamical variable. Using the URC empiric law as an {\\it ansatz} for this velocity immediately leads to expressions for the state variables that are (under suitable approximations) consistent with the thermodynamics of the non-relativistic Maxwell-Boltzmann ideal gas. From these expressions and bearing in mind numerical estimates of the empiric parameters appearing in the URC (``terminal'' rotation velocities and the ``core radius''), we obtain: (1) a constraint on the ratio of the particles mass to temperature for this gas, (2) the criterion of applicability of the Maxwell-Boltzmann distribution ({\\it i.e.} the non-degeneracy criterion) \\cite{Landau}, leading to a minimal bound of about $30$ to $60$ eV for the mass of the gas particles. Therefore, the assumption of a Maxwell-Boltzmann gas (SCDM or WDM) model for the galactic halo leads to an acceptable value for the particle's mass lying in the range $m > 0.5\\,\\hbox{keV}$. We provide in Table~1 a list of particle candidates that could be accomodated according to the criteria (1) and (2) above, namely: neutralino, photino, light gravitino, sterile neutrino, dilaton, axino, majoron, mirror neutrino and possibly standard massive neutrinos. As mentioned previously, this mass range is compatible with predictions of current work based on SCDM and WDM structure formation models. We find it interesting to remark that barions and electrons comply with the criterion (2) above, but (1) would imply gas temperatures of the order of $10^{3}-10^{6}$ K. A gas of barions or electrons at such temperatures would certainly not be ``dark''. HDM or WDM models based on less massive particles, like the electron neutrino, remain outside the scope of the present work, since these particles might require assuming a fuly relativistic Maxwell-Boltzmann gas or a degenerate gas (possibly relativistic) complying with a Fermi-Dirac or Bose-Einstein statistics. The axion, as well as other non-thermal relique sources, are also outside the scope of this paper and their study requires a different approach. The paper is composed as follows. In the next section we present the field equations for a static sphericaly symmetric spacetime with a perfect fluid source. We provide in section 3 a review of the thermodynamics of an equilibrium Maxwell-Boltzmann gas. Then, in section 4, we re-write the field equations in terms of the orbital velocity of stable circular geodesics and then assume for this velocity the empiric {\\it ansatz} given by the URC. This leads to forms of the state variables that will be compatible, under suitable series approximations, with the thermodynamics of the Maxwell-Boltzmann gas. Putting all these results together, we discuss in the last section the possible ranges for the mass of the particles of the dark halo gas and suggest future lines of research. \\setcounter{equation}{0} ", + "conclusions": "So far we have found a reasonable approximation for galactic dark matter to be described by a self gravitating Maxwell-Boltzmann gas, under the assumption of the empiric rotation velocity law (\\ref{eq:v}). The following consistency relations emerge from equations (\\ref{eq:n11}), (\\ref{eq:n22}) and (\\ref{eq:consist}) \\begin{equation} n_c \\ \\approx \\ \\frac{3\\,v_0^2}{4\\pi G\\,m\\,a^2\\,r_{\\hbox{\\tiny{opt}}}^2} ,\\qquad T_c \\ \\approx \\ \\frac{m\\,v_0^2}{3\\,k_{_B}} \\label{eq:nTc} \\end{equation} \\noindent hence, bearing in mind that $n\\leq n_c$ and $T\\approx T_c$, the condition (\\ref{eq:landau}) for the validity of the MB distribution can be written as \\begin{equation} \\frac{n\\,\\hbar^3}{(m\\,k_{_B}\\,T)^{3/2}} \\ \\leq \\ \\frac{n_c\\,\\hbar^3} {(m\\,k_{_B}\\,T_c)^{3/2}} \\ \\ll \\ 1, \\label{landau2} \\end{equation} \\noindent Inserting (\\ref{eq:nTc}) into (\\ref{landau2}) yields the condition \\begin{equation} m \\ \\gg \\ \\left[\\frac{3^{5/2}\\,\\hbar^3}{4\\pi G\\,a^2\\,r_{\\hbox{\\tiny{opt}} }^2\\,v_0}\\right]^{1/4}, \\label{aplicab} \\end{equation} \\noindent a criteria of aplicability of the MB distribution (non-degeneracy) that is entirely given in terms of $m$, the fundamental constants $G,\\,\\hbar $ and the empiric parameters $v_0$ and $b \\equiv a\\,r_{\\hbox{\\tiny{opt}}}$ (the ``terminal'' rotation velocity and the ``core radius''). For dark matter dominated galaxies (spiral and low surface brightness (LSB)) \\cite{urc_1} these parameters have a small variation range: $r_{\\hbox{\\tiny{opt}}} \\approx 15$\\,\\hbox{kpc} , $1$ kpc \\, $\\leq b \\leq$ \\, 5 kpc and $125\\, \\hbox{km/sec} \\leq v_0 \\leq 250\\, \\hbox{km/sec}$, the constraint (\\ref{aplicab}) does provide a tight estimate of the minimal value for the mass of the particles under the assumption that these particles form a self gravitating ideal dark gas complying with MB statistics. As shown in Figure 2, this minimal value lies between $30$ and $60$ eV, thus implying that appropriate particle candidates must have a much larger mass than this range of values. This minimal bound excludes, for instance, light mass particles such as the electron neutrino ($m_{\\nu_e} < 2.2$ eV) or the axion ($m_A \\approx 10^{-5}$ eV). The currently accepted estimations of cosmological bounds on the sum of masses for the three active neutrino species is about $24$ eV \\cite{PDG}, a value that would apparently rule out all neutrino flavours. However, recent estimations of these cosmological bounds have raised this sum to about $1$ keV \\cite{sn2}, hence more massive neutrinos could also be accomodated as dark matter particle candidates. Estimates of masses of various particle candidates are displayed in Table~1. Since $T\\approx T_{c}$, the consistency condition (\\ref{eq:consist}) provides the following constraint on the temperature and particles mass of the dark gas \\begin{equation} \\frac{m}{T_{c}}\\ =\\ \\frac{3\\,k_{_{B}}}{v_{0}^{2}}\\ \\approx 0.4\\times 10^{3}\\, \\frac{\\hbox{eV}}{\\hbox{K}} \\ , \\label{eq:consist2} \\end{equation} \\noindent where we have taken $v_{0}=200\\,\\hbox{km/sec}$ \\footnote{The variation of $v_0$ in the observed ranges for spiral galaxies does not alter significanly the numerical value in the rhs of Eq. (\\ref{eq:consist2}).}. Considering in (\\ref{eq:consist2}) the minimal mass range that follows from (\\ref{aplicab}), we would obtain gas temperatures consistent with the assumed typical temperatures of relic gases: $T_{c}\\approx 2\\,\\,\\hbox{to}\\,\\,4\\,\\hbox{K}$. However, as long as we do not have more information on the interaction and physical properties of various particle candidates, we cannot rule out a given large mass value on the grounds that the corresponding gas temperature could be too high. However, if we assume that the ideal dark gas is made of electrons or barions, so that $m=m_{p}$ or $m=m_{e}$, then condition (\\ref{aplicab}) for applicability of the MB distribution is certainly satisfied, but (\\ref{eq:consist2}) implies a temperature of the order of $T_{c}\\approx 10^{3}$ K for electrons and $T_{c}\\approx 10^{6}$ K for barions ! Obviuosly, barions or electrons at such a high temperatures would certainly not remain unobservably ``dark''. We can rule them out, but we cannot rule out more massive particles (in the range of 1-100 GeV's) characterized by weak interaction even if the gas temperature is in the range of $10^8 - 10^9$ K. Figure 3 illustrates, for various particle candidates, the relation between $T_c$ and $m$ contained in (\\ref{eq:consist2}). The main novelty of the present paper is the fact that it is based on a general relativistic hydrodynamics, as opposed to numerical simulations \\cite{nbody_1}-\\cite{nbody_3}, newtonian or Kinetic Theory perturbative approaches (see \\cite{scdm_1}-\\cite{wdm_6}). Finaly, the fact that under the assumption of MB distribution, we have obtained a minimal mass on the range $30 \\,\\hbox{eV} \\leq m \\leq 60 \\,\\hbox{eV}$ that seems to discriminate against thermal relique gases composed by lighter particles (electron neutrino, etc) coincides with the fact that these particle candidates tend to be ruled out because of their inability to produce sufficient matter clustering \\cite{tbook_1}, \\cite{tbook_2}. In spite of these arguments, if either of these particles constitute a self gravitating gases accounting for a galactic halo it would be inconsistent to model such a gas as SCDM in the context of a classical ideal gas complying with MB statistics. It would be necessary to examine these cases as either HDM or WDM, by using either a relativistic MB distribution (very light particles can be relativistic even at low temperatures) and/or a distribution that takes into account (depending on the particle) Fermi-Dirac or Bose-Einstein statistics. Non-thermal axions are very light particles ($m \\sim 10^{-5}$ eV), however this type of relique source cannot be treated as a Maxwell-Boltzmann gas, thus the lower mass limit that we have obtained does not apply. This and other non-thermal sources \\cite{sfe_1, sfe_2} require a wholy different approach. These studies will be undertaken in future papers. \\begin{table}[tbp] \\begin{center} \\begin{tabular}{|c|c|c|} \\hline {SCDM/WDM} & {mass in keV} & {References} \\\\ {Light Candidates } & {} & {} \\\\ \\hline\\hline {Light Gravitino\\hfill} & {$\\sim 0.5$} & {\\ \\cite{g1} } \\\\ {\\hfill} & {$\\sim 0.75 - 1.5$ } & {\\cite{g2}, \\cite{g3} } \\\\ \\hline {\\hfill} & {$\\sim 2.6 - 5$} & {\\cite{ssn1}} \\\\ {Sterile Neutrino\\hfill} & {$< 40$ } & {\\cite{ssn2} } \\\\ {\\hfill} & {$1 - 100$ } & {\\cite{ssn3} } \\\\ \\hline {Standard Neutrinos\\hfill} & {$\\sim 1$} & {\\cite{sn1}, \\cite{sn2} } \\\\ \\hline {Dilaton\\hfill} & {$\\sim 0.5$} & {\\cite{d1} } \\\\ \\hline {Light Axino\\hfill} & {$\\sim 100$} & {\\cite{a1} } \\\\ \\hline {Majoron\\hfill} & {$\\sim 1$} & {\\cite{m1}, \\cite{m2}, \\cite{m3} } \\\\ \\hline {Mirror Neutrinos\\hfill} & {$\\sim 1$} & {\\cite{mn1}, \\cite{mn2} } \\\\ \\hline \\hline {CDM} & {mass in GeV} & {References} \\\\ {Heavy Candidates } & {} & {} \\\\ \\hline\\hline {Neutralino\\hfill} & {$> 32.3$} & {\\ \\cite{abreu} } \\\\ {\\hfill} & {$> 46$ } & {\\cite{ellis} } \\\\ \\hline {Axino\\hfill} & {$\\sim 10$} & {\\cite{covi}, \\cite{leszek}} \\\\ \\hline {Gravitino\\hfill} & {${\\buildrel <\\over \\sim }\\ 100$} & {\\cite{kawasaki} } \\\\ \\hline \\end{tabular} \\end{center} \\caption{{\\em Particle candidates for a MB Dark Matter gas.}} \\end{table}" + }, + "0112/astro-ph0112235_arXiv.txt": { + "abstract": "This review presents a perspective on recent advances in understanding neutral ISM structure in external galaxies. \\hi\\ is a fundamental probe of galactic baryonic material, and its structure and distribution offer vital signatures of dynamical and evolutionary processes that drive star formation and galaxy evolution. New, high-resolution \\hi\\ data cubes for external galaxies now reveal the features and topology of the entire neutral ISM, which here are considered on scales of 10 -- 1000 pc. I focus on the two principal candidates for \\hi\\ structuring, mechanical feedback from massive stars and turbulence; other mechanisms are also considered, especially with respect to supergiant shells. While confirmation for both mechanical feedback and turbulent processes exists, it remains unclear how these mechanisms yield the global, steady-state, scale-free \\hi\\ properties that are observed. Understanding the formation of filamentary structure may be key in resolving these puzzles. New \\hi\\ surveys of nearby galaxies, combined with further theoretical studies, promise continuing important advances. ", + "introduction": "The distribution of neutral hydrogen in galaxies is an essential tracer of structure in the interstellar medium (ISM), and hence, of dynamical and evolutionary processes that drive structure formation. The \\hi\\ distribution itself, which often dominates the gas mass in galaxies, is one of our principal probes of galactic baryonic material. In particular, \\hi\\ mapping of late-type galaxies is well-known to reveal a gas distribution that can extend many times the characteristic optical or stellar radius of the galaxy; vivid examples are NGC 4449 (Hunter {\\etal}1998) and NGC 6822 (de Blok \\& Walter 2000), among many others. These galaxies show massive extended disk structure and tidal features that are unseen in other galactic components. I will present here an extragalactic, ``user's'' perspective on the role of \\hi\\ structure as a probe of interstellar processes relevant to galaxy evolution. The \\hi\\ structure bears directly on phenomena that are fundamental to galaxy evolution and star formation: mechanical feedback from massive stars and their supernovae (SNe); cloud formation and star formation; porosity of the cool ISM, especially to ionizing radiation; ISM phase balance and physics of phase interface regions; interstellar mixing and chemical enrichment. These processes are manifested in the neutral ISM in varying ways: superbubbles and shells presumably result from massive star mechanical feedback; fractal structure has been associated with turbulence; filaments result from various processes including feedback and magnetic activity; clouds result from gravitational effects and influence from other processes; tidal features and spiral structure are dominated by gravitation. Structure on scales of 10 -- 1000 pc is thought to be dominated by two principal processes: mechanical feedback from massive stars and turbulence, which are discussed in turn below. At larger spatial scales, gravitational processes will dominate, and these will not be considered here. The smallest spatial scales are best studied in the Galaxy and are discussed in these proceedings by Faison. ", + "conclusions": "The two primary candidate mechanisms for structuring the neutral ISM in galaxies, mechanical feedback from massive stars and turbulence, are both clearly significant effects. Signatures of both processes are confirmed in the \\hi\\ properties of external galaxies, and global characterizations of the ISM in terms of these processes are available. However, the specifics of the physical processes by which these mechanisms lead to the steady-state, observed global ISM properties are lacking. Neither mechanical feedback nor turbulence can exclusively explain the \\hi\\ topology. Given the rapid, localized turbulent dissipation process, a global role for turbulent structuring remains to be confirmed. Further studies making use of \\hi\\ and other datasets are therefore essential to clarify the circumstances under which these respective mechanisms dominate, as well as to identify additional relevant mechanisms, e.g., HVC impacts, gravitational and magnetic effects, and gas instabilities. This work must ultimately lead to an integrated view of structuring and dynamics in the multi-phase ISM and the corresponding consequences for star formation and galaxy evolution." + }, + "0112/astro-ph0112439_arXiv.txt": { + "abstract": "Five observations of the hard X-ray spectrum of Her X-1 from \\xte\\ show that the $\\sim \\! 41$ keV energy of the cyclotron scattering line is constant within statistics of a few percent per observation. The overall spectral shape, on the other hand, varies somewhat, with an RMS of 2\\%. If the 41 keV feature truly originates as cyclotron resonance scattering in an unchanging 3 $\\times$ 10$^{12}$ Gauss dipole field not far above the neutron star surface, these observations constrain the average height of scattering to within a range of 180 meters. This is consistent with models which put the radiating structure within meters of the surface of the neutron star. In other pulsars observed line centroid changes have been correlated with luminosity changes, and if interpreted as variations of the height at which scattering takes place, many hundreds of meters are required. These \\xte\\ data, which sample nearly a factor of two in unabsorbed luminosity, are in conflict with a particular model for such an extended radiating structure. Comparison with other observations over many years indicates strongly that the centroid energy of this absorption line has increased some time between 1991 and 1993 by 23\\%, from 34 keV to 41 keV. Moreover, the cutoff energy of the spectral continuum increased at the same time from 16 keV to 20 keV, which is, within the statistical error of 5\\%, in direct proportion to the centroid. This may be a sign that both these characteristics of the spectrum are controlled in the same way by the magnetic field strength in the region of scattering. ", + "introduction": "The accreting X-ray pulsar, Hercules X-1, emits a spectrum with a localized depression at roughly 40 keV which is generally interpreted as an absorption line resulting from scattering of photons on electrons whose motions are constrained in one spatial dimension to transitions between Landau levels in a teraGauss magnetic field at the neutron star's polar cap. Commonly referred to as a cyclotron line, this feature was discovered in 1977 (Tr\\\"{u}mper et al. 1978), although first identified phenomenologically as an emission feature at somewhat higher energy. This feature in Her X-1 has since been observed repeatedly (e.g. Gruber et al. 1980, Voges 1984, Tueller et al. 1984, Soong et al. 1988, Mihara 1995, Kunz 1995, Dal Fiume et al. 1998). Similar features have also been identified in the hard X-ray spectra of many of the accreting X-ray pulsars (e.g. Mihara 1995, Heindl et al. 1999, Dal Fiume et al. 2000). Cyclotron lines afford a diagnostic of plasmas under extreme conditions and can provide a key to the structure of the magnetosphere and the mass flow. Mihara et al. (1998) have reported a study of repeated observations with \\ginga\\ of cyclotron lines in a number of X-ray pulsars. In a few instances, although not with Her X-1, they found a correlation with X-ray luminosity which was consistent with the cyclotron scattering taking place from hundreds to many hundreds of meters above the neutron star surface, at the peak of an ``accretion mound\" structure proposed by Burnard, Arons \\& Klein (BAK, 1991). BAK predict a strong dependence of this height on accretion flow. Thus in this model variability of the cyclotron line energy can also be an important diagnostic of the scattering region and accretion process. Studies of line formation (M\\'{e}sz\\'{a}ros \\& Nagel 1985, Isenberg, Lamb \\& Wang 1998, Araya-Gochez \\& Harding 2000) have indicated the possibility of line profiles which deviate from the simple Gaussian and Lorentzian profiles which have been used for analysis to date. Such deviations are evident as particularly broad lines (Heindl 1999, Orlandini 1999). However, the line of Her X-1 is narrower, and measurement of structure beyond the first and second moments (centroid and width) requires detector resolution higher than the $\\delta E/E$ of a few obtainable with inorganic scintillators, as well as very large collecting area and observing time. Using data from five observations over two years with the large-area instruments aboard Rossi X-Ray Timing Explorer (\\xte ) we have performed detailed spectroscopy of the spectral continuum, the cyclotron line and the fluorescent iron line of Her X-1. We find only minor variability between these observations. Exercising careful data selection, we have compared the results with those from earlier measurements between 1977 and 1994. For this purpose the \\heao\\ data were freshly analyzed with modern software and spectral models. To a level of perhaps 10\\% the pre-1991 spectral shapes are consistent with each other, as are the spectra from 1994 onwards. Between 1991 and 1994 there strongly appears to have been a change of the cyclotron line centroid and the cutoff energy of the continuum. This is in marked contrast to the power law index and exponential fold energy, for which the average values remained constant to within a few percent. In \\S 2 we describe instrumentation and observations, in \\S 3 analysis and results, and in \\S 4 we briefly discuss possible interpretations. Preliminary results from a portion of this \\xte\\ data set have been reported earlier (Gruber et al. 1997, 1998, 1999). ", + "conclusions": "We have shown convincing evidence for a 23\\% historical change in the cyclotron line energy in the spectrum of Her X-1 through comparison of \\xte\\ observations with earlier observations. For this purpose archival \\heao\\ spectra were re-analyzed, using modern tools and spectral forms for consistency. On the other hand, the five \\xte\\ observations, which span about two years, show no change, with a limit of about three percent. We also demonstrate a historical change of continuum cutoff energy which is closely proportional to the cyclotron centroid energy. This is the first good evidence to connect the cutoff energy, which is common to accreting pulsars, to the magnetic field. These \\xte\\ data do not support a particular accretion mound scenario." + }, + "0112/astro-ph0112113_arXiv.txt": { + "abstract": "Evolutionary models have been calculated for Pop II stars of 0.5 to 1.0$\\Msol$ from the pre-\\MS{} to the lower part of the giant branch. Rosseland opacities and radiative accelerations were calculated taking into account the concentration variations of 28 chemical species, including all species contributing to Rosseland opacities in the OPAL tables. The effects of radiative accelerations, thermal diffusion and gravitational settling are included. While models were calculated both for $Z=0.00017 $ and $0.0017$, we concentrate on models with $Z=0.00017 $ in this paper. These are the first Pop II models calculated taking radiative acceleration into account. It is shown that, at least in a 0.8 \\Msol{} star, it is a better approximation not to let \\Fe{} diffuse than to calculate its gravitational settling without including the effects of \\gr(Fe). In the absence of any turbulence outside of convection zones, the effects of atomic diffusion are large mainly for stars more massive than 0.7$\\Msol$. Overabundances are expected in \\emph{some} stars with $\\teff{} \\ge 6000 $\\Kelvin{}. Most chemical species heavier than CNO are affected. At 12 Gyr, overabundance factors may reach 10 in some cases (e.g. for Al or Ni) while others are limited to 3 (e.g. for Fe). The calculated surface abundances are compared to recent observations of abundances in globular clusters as well as to observations of Li in halo stars. It is shown that, as in the case of Pop I stars, additional turbulence appears to be present. Series of models with different assumptions about the strength of turbulence were then calculated. One series minimizes the spread on the Li plateau while another was chosen with turbulence similar to that present in AmFm stars of Pop I. Even when turbulence is adjusted to minimize the reduction of Li abundance, there remains a reduction by a factor of at least 1.6 from the original Li abundance. Independent of the degree of turbulence in the outer regions, gravitational settling of He in the central region reduces the lifetime of Pop II stars by 4 to 7 \\% depending on the criterion used. The effect on the age of the oldest clusters is discussed in an accompanying paper. Just as in Pop I stars where only a fraction of stars, such as AmFm stars, have abundance anomalies, one should look for the possibility of abundance anomalies of metals in some Pop II turnoff stars, and not necessarily in all. Expected abundance anomalies are calculated for 28 species and compared to observations of M92 as well as to Li observations in halo field stars. ", + "introduction": "\\label{sec:context} Helioseismology has confirmed the importance of gravitational settling in the Sun's external regions (\\citealt{GuzikCo92}; \\citealt{GuzikCo93}; \\citealt{ChristensenDalsgaardPrTh93}; \\citealt{Proffitt94}; \\citealt{BahcallPiWa95}; \\citealt{GuentherKiDe96}; \\citealt{RichardVaChetal96}; \\citealt{BrunTuZa99}). Turnoff stars in globular clusters are only slightly less massive than the Sun and have convection zones that tend to be somewhat shallower. In solar type stars, radiative accelerations have been shown to become equal to gravity for some metals around the end of the main sequence lives \\citep{TurcotteRiMietal98}. The question of the atomic diffusion of metals in Pop II stars then naturally arises. It is currently of special interest because large telescopes are now making possible the determination of the abundance of metals in the turnoff stars of globular clusters. In this paper evolutionary models that take into account the diffusion of He, LiBeB and metals in Pop II stars are presented for the first time. Surface abundances may then be used as additional constraints in the determination of the age of globular clusters and of the Universe (\\citealt{VandenBergRiMietal2001}, hereafter Paper II). Given their old age, Pop II stars are those where the slow effect of atomic diffusion has the largest chance to play a role on the evolutionary properties. Previously published evolutionary models of Pop II stars have included some of the effects of diffusion (\\citealt{DeliyannisDeKa90}; \\citealt{ProffittMi91}; \\citealt{ProffittVa91}; \\citealt{SalarisGrWe2000}; \\citealt{StringfellowBoNoetal83}) but never included the effects of the diffusion of metals with their radiative accelerations self consistently. Determining constraints on stellar hydrodynamics from abundance observations requires knowing the original chemical composition of the star. In Pop II stars there are larger variations in original abundances than in Pop I. For that reason it is essential to use globular clusters for abundance determinations of metals since only then does one have a handle on the original abundances from observations of cluster giants. However LiBeB have an origin partly different from that of most metals. Coupled with their sensitivity to low temperature nuclear burning, their abundance determination in halo stars provides useful constraints on hydrodynamics even if determinations in cluster stars would be preferable. \\footnote{Only in the case of NGC 6397 \\citep{MolaroPa94} and M92 \\citep{BoesgaardDeKietal99} have Li abundances been determined for turnoff stars, and these data are much less precise than those obtained for halo field stars. The M92 data are discussed later in the paper. } The observation by \\citet{SpiteSp82} of a plateau in the Li concentration over a relatively large \\teff{} interval of Pop II stars has now been confirmed by many observations. Furthermore, the Li concentration is constant while that of Fe varies by more than a factor of 100 \\citep{Cayrel98} from $[\\Fe/\\H] = - 3.7$ to $-1.5$. This shows the primordial origin of Li. Its preservation for such a time interval over such a wide \\teff{} interval seriously challenges our understanding of convection and other potential mixing processes in those stars \\citep{MichaudFoBe84}. If there were no mixing process outside of convection zones, the surface Li abundance would vary with \\teff{}: at small \\teff{} because of nuclear reactions ($^7\\Li (p,\\alpha)^4\\He$) and at large \\teff{} because of gravitational settling. Extending convection zones by a simple turbulence model does not solve the problem since, if the extension is sufficient to reduce Li settling enough in the hotter stars, it causes excessive Li destruction in the cooler stars of the plateau \\citep{ProffittMi91}. Efforts were also made to link such an extension to differential rotation (\\citealt{Vauclair88} or, as parametrized in the Yale models, \\citealt{PinsonneaultWaStetal99}) but the small apparent dispersion in the plateau makes this model unlikely. \\citet{RyanNoBe99} even claim that the destruction may not be by more than 0.1 dex. Several groups (\\citealt{CayrelSpSpetal99}, \\citealt{HobbsTh94, HobbsTh97}, \\citealt{SmithLaNi93}) have also observed $^6\\Li$ in halo stars. Since $^6\\Li$ is destroyed at a smaller \\teff{} than $^7\\Li$, its survival in old stars implies a strict and small upper limit on the amount of mixing in those stars. The stars where a detection has been made are all concentrated close to the turnoff \\citep{CayrelSpSpetal99}. While a mechanism has been suggested to produce $^6\\Li{}$ in the Sun \\citep{RamatyTaThetal2000}, it is not expected to significantly affect atmospheric values so that a star in which $^6\\Li{}$ is seen in the atmosphere may not have destroyed a significant fraction of its original $^6\\Li$. \\citet{BoesgaardDeKietal99} determined Be abundances in halo stars as well as a linear correlation with the Fe abundance suggesting that Be has not been destroyed in those stars. The chemical composition of globular clusters is attracting more attention as large telescopes make the determination of the surface chemical composition of turnoff stars possible. The determination of the abundance of metals makes it possible to put additional constraints on the hydrodynamics of Pop II stars. Some observers have reported factors of 2 differences between Fe abundances in red giants and subgiants \\citep{KingStBoetal98}. However \\citet{RamirezCoBuetal2001} and \\citet{GrattonBoBretal2001} find no difference between the \\Fe{} abundances in red giants and turnoff stars. \\citet{TheveninChdeetal2001} find the same relative abundances as in the Sun in the turnoff stars of NGC$\\,$6397. Unexplained anticorrelations between the abundances of O and Na have been observed by \\citet{GrattonBoBretal2001}. In this paper, the surface abundances to be expected in Pop II stars are calculated under different assumptions for the internal stellar hydrodynamics. The first and simplest assumption is that there is no macroscopic motion outside of convection zones. There remains only atomic diffusion, including the effects of gravitational settling, thermal diffusion and radiative accelerations, as a transport process outside of convection zones. On the main sequence, such models have been shown to lead to larger abundance anomalies than are observed \\citep{TurcotteRiMi98} but also to the appearance of an additional convection zone caused by the accumulation of iron peak elements \\citep{RichardMiRi2001}. A relatively simple parametrization of turbulence, corresponding to an extension of those Fe convection zones by a factor of about 5 in mass, was shown to lead to a simple explanation of the AmFm phenomenon \\citep{RicherMiTu2000}. A similar parametrization is used here in Pop II stars and leads to our second series of models. Finally, we use an additional parametrization that is chosen in order to minimize the reduction of the surface Li concentration, so as to provide the best representation of the Spite plateau. \\citet{ProffittMi91} introduced a similar parametrization of turbulence to compete with He and Li settling (other parametrizations of turbulence were introduced for the Sun by \\citealt{RichardVaChetal96} and \\citealt{BrunTuZa99}). \\citet{Basu97} has shown that weak turbulence below the solar convection zone also improves agreement with the solar pulsation spectrum. For comparison purposes, series of models are also calculated without diffusion and one model is calculated with gravitational settling but without \\gr{} (see subsubsection~\\ref{sec:structure}). ", + "conclusions": "\\label{sec:conlusion} It has been clearly shown that, contrary to the belief expressed in many papers and recently by \\citet{ChaboyerFeNeetal2001}, atomic diffusion does not necessarily lead to underabundances of metals in Pop II stars. Differential radiative accelerations lead to overabundances of Fe and some other chemical elements in some turnoff stars. Consider the evolution of Pop II stars with no tubulence. As one may see by considering the solid curves in Figures~\\ref{fig:ab_teff12} and \\ref{fig:ab_teff13.5}, generalized underabundances by 0.1 dex are expected in the \\teff{} interval from 4600 to 5500 \\Kelvin{}. Between 5500 and 6000 \\Kelvin{}, the underabundances are still generalized and increase to 0.3 dex for some species such as \\Ox{}. In 12 Gyr turnoff stars however ($\\teff{} \\geq 6000$ \\Kelvin{}), overabundances by a factor of up to 10 are possible (e.g. Al and Ni). All calculated species heavier than Na may have overabundances. At a given \\teff{}, variations are expected from star to star. At 13.5 Gyr, similar but smaller anomalies are expected. The overabundances are sensitive to any left over turbulence below the convection zone. In this paper, the evolution of stars with $Z = 0.00017$ has been described both with and without turbulence. A 0.1 dex underabundance of metals in turnoff stars as compared to giants has been shown to be the smallest anomaly to be expected (section \\ref{sec:ab_teff}). Star-- to-- star variations were seen to be possible around the turnoff, if turbulence is small enough. Observations (see section \\ref{sec:observations}) suggest the presence of abundance variations similar to those expected. The comparison to observations is, however, sensitive to the \\teff{} scale. As \\citet{KingStBoetal98} concluded at the end of their section 4.2, higher quality data is probably required to establish the reality of Fe abundance variations within M92. The accurate determination of the abundance of more species is also needed. This may well have implications not only for intrinsic abundance variations in clusters but for internal stellar structure. The effect of varying $Z$ on the evolution of Pop II stars will be investigated in a forthcoming paper, (Paper III), where comparisons to higher $Z$ clusters will be undertaken. Increasing $Z$ in Pop II stars will be shown to reduce considerably the expected abundance anomalies. Note also that Paper II shows that the present models for $[\\Fe/\\H]_0 = -1.31$ accurately reproduce the CMD locations of local Population II subdwarfs having Hipparcos parallaxes and metallicities within +/$-$ 0.2 dex of $[\\Fe/\\H] = -1.3$. In a number of clusters with higher $Z$ than M92, observations suggest that only small variations, if any, are present in turnoff stars (see for instance \\citealt{RamirezCoBuetal2001}; \\citealt{GrattonBoBretal2001} and \\citealt{TheveninChdeetal2001}). Furthermore in field halo stars, the Li abundance puts strict constraints on any chemical separation. In the companion paper (Paper II) we therefore took the cautious approach to use mainly evolutionary models that minimize the effect of atomic diffusion. It has been shown that the use, in complete stellar evolution models, of a relatively simple parametrization of turbulent transport leads to Li surface abundances compatible with the Li plateau observed in field halo stars (with a 0.17 dex reduction from the original Li abundance) and small variations in the surface abundances of metals (a 0.1 dex reduction of the metal abundance in turnoff stars as compared to that in giants in clusters with $Z = 0.00017$). At the same time, the gravitational settling of He leads to a reduction in the age of globular clusters by some 10\\,\\% (see Paper II). However simple the parametrization of turbulent transport, it is not understood from first principles. The high level of constancy of Li abundance as a function of \\teff{} requires that turbulence mixes to very nearly the same $T$ throughout the star evolution and in stars covering a mass interval of approximately 0.6 to 0.8 \\Msol{}. As already noted by \\citet{MichaudFoBe84}, this is not expected in standard stellar models. No convincing hydrodynamic model has been proposed that explains this property. Pop II stars appear to tell us that this is the case, however. Mass loss is another physical process that could compete with atomic diffusion and maintain a constant value of Li as a function of \\teff{} \\citep{VauclairCh95}. Whether, in the absence of turbulent transport, it could be made consistent with the observations of metals in globular cluster turnoff stars is a question that requires further calculations. These may lead to observational tests of the relative importance of mass loss and turbulence in these objects. The number of chemical species that are now included in these calculations and that can be observed makes such tests possible." + }, + "0112/astro-ph0112325_arXiv.txt": { + "abstract": "Neutral hydrogen VLA D-array observations of the dwarf irregular galaxy HoII, a prototype galaxy for studies of shell formation, are presented. These were extracted from the multi-configuration dataset of Puche et al.\\ (\\markcite{PWBR92}1992). \\ion{H}{1} is detected to radii over 16\\arcmin\\ or $4R_{25}$, almost a factor of two better than previous studies. The total \\ion{H}{1} mass $M_{\\mbox{\\scriptsize HI}}=6.44\\times10^8$~\\msun. The integrated \\ion{H}{1} map has a comet-like appearance, with a large but faint component extending to the northwest and the \\ion{H}{1} appearing compressed on the opposite side. This suggests that HoII is affected by ram pressure from an intragroup medium (IGM). The velocity field shows a clear rotating disk pattern and a rotation curve corrected for asymmetric drift was derived. However, the gas at large radii may not be in equilibrium. Puche et al.\\ (\\markcite{PWBR92}1992) multi-configuration data were also reanalyzed and it is shown that they overestimated their fluxes by over 20\\%. The rotation curve derived for HoII is well defined for $r\\lesssim10$~kpc. For $10\\lesssim r\\lesssim18$~kpc, however, velocities are only defined on the approaching side, such that this part of the rotation curve should be used with caution. An analysis of the mass distribution, using the whole extent of this rotation curve, yields a total mass of $6.3\\times10^9$~\\msun, of which $\\approx80$\\% is dark. Similarly to what is seen in many dwarfs, there is more luminous mass in \\ion{H}{1} than in stars. One peculiarity, however, is that luminous matter dominates within the optical body of the galaxy and dark matter only in the outer parts, analogous to what is seen in massive spirals rather than dwarfs. HoII lies northeast of the M81 group's core, along with Kar~52 (M81~Dwarf~A) and UGC~4483. No signs of interaction are observed, however, and it is argued that HoII is part of the NGC~2403 subgroup, infalling towards M81. A case is made for ram pressure stripping and an IGM in the M81 group. Stripping of the outer parts of the disk would require an IGM density $n_{\\mbox{\\tiny IGM}}\\gtrsim4.0\\times10^{-6}$~atoms~cm$^{-3}$ at the location of HoII. This corresponds to $\\sim$1\\% of the virial mass of the group uniformly distributed over a volume just enclosing HoII and it is consistent with the known X-ray properties of small groups. The \\ion{H}{1} tail is consistent with additional turbulent viscous stripping and evaporation, at least for low IGM temperatures. It is argued that existing observations of HoII do not support self-propagating star formation scenarios, whereby the \\ion{H}{1} holes and shells are created by supernova explosions and stellar winds. Many \\ion{H}{1} holes are located in low surface density regions of the disk, where no star formation is expected or observed. Alternative mechanisms are discussed and it is suggested that ram pressure can help. Ram pressure has the capacity to enlarge preexisting holes and lower their creation energies, helping to bridge the gap between the observed star formation rate and that required to create the holes. ", + "introduction": "} \\nopagebreak Spiral galaxies have a complex interstellar medium (ISM) pierced by numerous cavities. These were first identified in the Galaxy (Heiles \\markcite{h79}1979, \\markcite{h84}1984) and M31 (Brinks \\markcite{b81}1981; Brinks \\& Bajaja \\markcite{bb86}1986), but similar structures were quickly discovered in other nearby galaxies (e.g.\\ Deul \\& den Hartog \\markcite{dh90}1990 for M33; Kamphuis \\markcite{k93}1993 for M101 and NGC~6946). The majority of these cavities are generally thought to arise from the combined effects of supernova explosions (SNe) and stellar winds (see Tenorio-Tagle \\& Bodenheimer \\markcite{tb88}1988 and van der Hulst \\markcite{h96}1996 for reviews), consistent with a three-phase picture of the ISM (Cox \\& Smith \\markcite{cs74}1974; McKee \\& Ostriker \\markcite{mo77}1977). In late-type spirals and dwarfs, cavities are long-lived due to a combination of low shear, a shallow gravitational potential, the absence of spiral density waves, and large scaleheights. One of the best studied dwarf irregulars is without doubt HoII, in the nearby M81 group of galaxies. Its basic properties are listed in Table~\\ref{ta:basic}. Puche et al.\\ (\\markcite{PWBR92}1992; hereafter \\markcite{PWBR92}PWBR92) cataloged and characterized over 50 \\ion{H}{1} holes and shells in HoII using multi-configuration observations from NRAO's Very Large Array\\footnote{The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under a cooperative agreement by Associated Universities, Inc.} (VLA), arguing for sequential star formation events in the disk. This picture has since been criticized (e.g.\\ Rhode et al.\\ \\markcite{rswr99}1999), but it was already clear then that some holes required extremely high creation energies (up to $2\\times10^{53}$~ergs). \\markcite{PWBR92}PWBR92 pointed out that HoII probably contains some dark matter, but they commented only briefly on its large-scale mass distribution, although it may be crucial to understand the evolution of the shells and constrain the amount of gas expelled from the galaxy (e.g.\\ Silich et al.\\ \\markcite{sfpt96}1996; Efstathiou \\markcite{e00}2000). \\markcite{PWBR92}PWBR92 detected a small \\ion{H}{1} cloud to the southwest, pointing to a significant amount of \\ion{H}{1} at large radii. This is essential to constrain the dark matter distribution, but also hints at a disturbed large-scale ISM or intergalactic medium (IGM). This is vindicated in our study, which shows that HoII has an extended comet-like \\ion{H}{1} morphology, possibly caused by ram pressure. \\placetable{ta:basic} In this paper, we will focus on HoII's large-scale \\ion{H}{1} distribution, looking at how it can help us understand its current properties and unravel its evolutionary history. In \\S~\\ref{sec:data}, we reanalyze and discuss \\markcite{PWBR92}PWBR92's data, using exclusively the low-resolution D-array observations. In \\S~\\ref{sec:mass}, we derive HoII's rotation curve and discuss its mass distribution. The likelihood of ram pressure and the properties of a putative IGM are discussed in \\S~\\ref{sec:ram+environment}. In \\S~\\ref{sec:shells}, the creation of shells and supershells is discussed in light of all available observations. We summarize our results and briefly conclude in \\S~\\ref{sec:conclusions}. ", + "conclusions": "} \\nopagebreak With low resolution but sensitive VLA observations, we have shown that the large-scale structure of the \\ion{H}{1} disk in HoII is comet-like. Since HoII does not appear to be interacting with its neighbors and may be infalling towards to core of the M81 group, it is probably undergoing ram pressure stripping from an IGM. About 1\\% of the virial mass of the group must be contained in an IGM if the outer parts of HoII are to be stripped, consistent with X-ray observations of small groups. The length and morphology of the \\ion{H}{1} tail observed point to the additional effects of turbulent viscous stripping and thermal evaporation. Ram pressure is interesting with regard to the creation of the \\ion{H}{1} holes in HoII. Since the expected traces of the implied SF are not observed, the creation energies of the holes through SNe and stellar winds must be systematically overestimated (or one must reject altogether this formation scenario). Ram pressure offers a mechanism to enlarge pre-existing holes, no matter how they were created, and thus lower their creation energies. A proof of the existence of a sufficiently dense IGM remains, however, the missing element to support this suggestion. While detecting a diffuse hot IGM of density $10^{-6}-10^{-5}$~atoms~cm$^{-3}$ would be very hard, especially in the M81 group given its large extent on the sky, it may be easier to look for other signatures of its presence. These include leading bow shocks (for galaxies moving supersonically through the IGM), gravitationally focused wakes, and the ram pressure stripped material itself (see Stevens et al.\\ \\markcite{sap99}1999). These may be detectable with the new generation of X-ray telescopes ({\\em XMM} and {\\em Chandra}) and could also be used to constrain the orbital structure (anisotropy) of the group (Merrifield \\markcite{m98}1998). HoII is a prime target for a search since it lies on the outskirt of a poor group with a presumably low IGM temperature. Contrary to expectations, the effects of ram pressure stripping are most easily visible in low surface brightness, cool, poor groups and clusters. In these, galaxies are able to retain a substantial fraction of their ISM, leading to prominent tails (see, e.g., Model~1b of Stevens et al.\\ \\markcite{sap99}1999) Deep \\ion{H}{1} observations of the region around HoII, Kar~52, and UGC~4483 are also crucial. They could reveal evidence of interactions that were missed by the present observations and which could also explain HoII's large-scale morphology. They could also reveal or rule out the presence of dense, cold material in the vicinity of HoII (e.g.\\ HVCs), which may be related to the creation of some of the holes. Observations along these lines are underway with the D configuration at the VLA." + }, + "0112/astro-ph0112055_arXiv.txt": { + "abstract": "The nonthermal filaments in the Galactic Center constitute one of the great mysteries of this region of the Galaxy. We summarise the observational data on these filaments and critically review the various theories which currently outnumber the observed filaments. We summarise out theory for the longest of these filaments, the Snake, and discuss the relevance of this model for the other filaments in the Galactic Center region. The physics involved in our model for the Snake involves much of the physics that has dominated the career of Professor Don Melrose. In particular, the diffusion of relativistic electrons in the Snake is determined from the theory of resonant scattering by Alfv\\'en waves. ", + "introduction": "We would like to thank the organisers of this Festschrift for Professor Don Melrose for the opportunity to contribute this paper, summarising our recent research \\cite{bicknell01c} on one of the mysterious filaments in the Galactic Center, known as the Snake. First, however, on this occasion, we would like to record a few personal notes based on our somewhat different interactions with Don Melrose. Geoff first met Don, when he took up his first appointment at Mt. Stromlo and Siding Spring Observatories (now the Research School of Astronomy and Astrophysics) at the Australian National University. Geoff had started working on extragalactic jets at that time and this clearly involved the physics of particle acceleration on which Don was an expert and he was not. When Geoff started at Mt. Stromlo, he contacted Don, and the result of the ensuing collaboration was a paper \\cite{bicknell82a}, that created a lot of interest and half of which he would still defend today. Which half? The half that clearly sets out the relationship between hydrodynamic turbulent input and its dissipation via particle acceleration at the high wave number end of a turbulent cascade, to which they both contributed. That collaboration with an outstanding scientist was one of the most memorable experiences of his career at that time and one that he always looks back upon with satisfaction. Jianke worked in the ARC Special Research Centre for Theoretical Astrophysics that Don headed from 1991 to 2000. Jianke was impressed with Don's vision for creating a career path for young theoretical astrophysicists and with his capacity for research and administration. Jianke also wrote a paper with Don on pulsar magnetospheres \\cite{li94a} and during the course of this work was impressed by Don's strongly focussed approach to Science, by his physical insight and by his insistence on understanding a problem fully at every stage. We both wish Don well in continuing his wide interests in Astrophysics, and in his new and challenging position as head of the School of Physics at Sydney University. We trust that he will continue to contribute to Astrophysics well into the future. The work that is the subject of this paper fortuitously involves physics to which Don has made significant and enduring contributions. Our work on the curious filament in the Galactic Center known as ``The Snake'' involves the following: \\begin{itemize} \\item Magnetic fields. \\item Particle acceleration. \\item Resonant scattering of relativistic particles and the diffusion of relativistic particles in the interstellar medium. \\item A strong connection with solar physics. \\end{itemize} Those of you who are familiar with Don's work will know that he has contributed significantly to all of these areas. The observations of the Snake also involved the outstanding Ph.D. thesis work of Sydney University student Andrew Gray and his advisors and colleagues Professor Lawrence Cram, Professor Ron Ekers, Dr. Miller Goss and Dr. Jenny Nicholls. The Snake was discovered through observations with the Molonglo Synthesis Telescope \\cite{gray91a}. It is therefore doubly appropriate to discuss this work on this occasion. In this paper, as well as summarising the work in our recent ApJ Letter, we have taken the opportunity to review some of the recent theoretical papers on the filaments, to expand some of the details in our own recent paper, to extend the treatment of relative timescales by incorporating a discussion of the important constraint imposed by synchrotron cooling, and to indicate where the physics discussed in this paper may relate to the Galactic Center filaments as a whole. ", + "conclusions": "We have summarised the observational situation and a number of theories for the curious filaments in the Galactic Center and have discussed at some length our own theory for the Snake. In the process of this brief review and summary of one theory, has anything been learned? The answer to this question is necessarily subjective and other workers in this field would answer this in entirely different ways. From a purely subjective viewpoint therefore, it seems to us that reconnection driven by some dynamical process together with electron diffusion and synchrotron cooling are the essential ingredients for a comprehensive theory of these filaments. There is good (circumstantial?) evidence for this: The sites of particle acceleration in the Arc are plausibly related to reconnection brought about by the interaction of molecular clouds with the magnetic field in Sagittarius~A, the increase in polarised emission at the crossing of two strands in G359.54+0.18 and the coincidence of peaks in radio emission at kinks in the Snake. It also seems to us that the sometimes bifurcated, sometimes multi-stranded, sometimes braided morphology of the filaments is indicative (or at least suggestive) of the topological rearrangement of magnetic field lines resulting from reconnection. We have quoted some work on this relating to twisted magnetic flux tubes and this has motivated the model we have advanced for the Snake. Our proposal for the production of kinks and reconnection through the rotation of the anchoring clouds stands or falls by the detection or non-detection of rapid rotation in molecular cloud/HII region cores which intersect the filaments. However, there are other ways in which magnetic flux tubes may interact to provide reconnection sites. Some recent work in a solar physics context involving colliding flux tubes (albeit twisted) is that by \\citeN{linton01a}. Whatever way reconnection is initiated, it seems that the interaction between molecular clouds and filaments is strongly related to the gas dynamics of the bar-driven accretion in the Galactic Center. Development of this theme seems to be an exciting and productive prospect and may illuminate the processes of accretion in galactic nuclei in general. Once electrons are accelerated at a given site they diffuse and cool, the latter mainly as a result of synchrotron losses. We have summarised two models that take diffusion into account, and it is interesting to note that the diffusion parameter in each case is consistent with the same level of turbulence and the $0.4 - 1 \\> \\rm mG$ strength of the magnetic field. For the Snake, we have argued that radiative losses are unimportant at the observed frequencies. However, other filaments that are older in terms of their cooling timescales would be expected to exhibit cooling features in their spectra. Therefore, it is unsurprising to see a variety of spectral index characteristics in the filaments. Cooling has been successfully incorporated into the diffusive model for the southern plume (connected to the Arc) and presumably we shall soon see diffusive plus cooling models for all of the NTFs. The issue of the particle spectrum resulting from the reconnection regions with or without associated shocks has received little attention to date. The theory of particle acceleration in shocks is well advanced; the theory of reconnection-induced particle acceleration less so although there has been some recent work in this area (eg. \\citeN{schopper99a}, \\citeN{birk01a}). A significant problem in the context of the filaments is what determines the parameters of the electron distribution, total energy density, minimum Lorentz factor etc. The strength of the magnetic field in the Galactic Center permeates all of the theoretical ideas that we have summarised. The two main contenders seem to be (1) A pervasive milli-Gauss field (2) Isolated instances of force-free fields. The helical field structure surrounding the Arc is good evidence for the latter and the existence of the filaments G358.85+0.47 and G359.85+0.39 parallel to the Galactic Plane tend to argue against the former. However, we are sure that this will continue to be a disputed point for some time and there are counterarguments -- such as the idea that the parallel filaments mark a change in direction of the Galactic Center magnetic field. Extrapolating our ideas on the Snake to other filaments, we attribute the predominance of filaments perpendicular to the plane to the lack of shear induced disruption for filaments in this direction. (This point arose in discussion with Professor Ron Ekers following the presentation at the Festschrift.) \\paragraph{Acknowledgements.} We are grateful to an anonymous referee for constructive comments and to Professors Ken Freeman and James Binney for useful discussions." + }, + "0112/astro-ph0112263_arXiv.txt": { + "abstract": "The ultraluminous broad absorption line quasar APM08279+5255 is one of the most luminous systems known. Here, we present an analysis of its nuclear CO(1-0) emission. Its extended distribution suggests that the gravitational lens in this system is highly elliptical, probably a highly inclined disk. The quasar core, however, lies in the vicinity of naked cusp, indicating that APM08279+5255 is truly the only odd-image gravitational lens. This source is the second system for which the gravitational lens can be used to study structure on sub-kpc scales in the molecular gas associated with the AGN host galaxy. The observations and lens model require CO distributed on a scale of $\\sim 400$ pc. Using this scale, we find that the molecular gas mass makes a significant, and perhaps dominant, contribution to the total mass within a couple hundred parsecs of the nucleus of APM08279+5255. ", + "introduction": "Identified serendipitously in a search for high latitude carbon stars, the $z=3.9$ broad absorption line quasar \\apm\\ is coincident with an IRAS source with a flux of 0.95Jy at 100$\\mu$m (Irwin et al. 1998). Observations with SCUBA reveal that \\apm\\ possesses a significant submillimetre flux of 75mJy at 850$\\mu$m (Lewis et al. 1998), implying a bolometric luminosity of $\\sim5\\times10^{15}L_\\odot$. Imaging reveals that \\apm\\ is not point-like, but rather is extended over a fraction of an arcsecond with a structure indicative of gravitational lensing (Irwin et al. 1998). The composite nature of \\apm\\ was confirmed in adaptive optics (AO) images obtained by Ledoux et al. (1998), with the system appearing as a pair of point-like images separated by 0.4 arcsec. Observations with NICMOS on the Hubble Space Telescope (Ibata et al. 1999) and AO images taken with the Keck telescope (Egami et al. 2000) also uncovered a fainter third component between the brighter two. Gravitational lens models derived from these observations suggest that the quasar continuum source has been magnified by $\\sim90$. Using IRAM, Downes et al. (1999) detected emission in CO(4-3) and CO(9-8), revealing the presence of warm circumnuclear gas in \\apm. Papadopoulos et al. (2001) were able to search for CO(1-0) and CO(2-1) in this system using the Very Large Array (VLA). Both were clearly detected associated with the quasar nucleus, as well as a more extended component located several arcsecs from the quasar images. Using locally established values of the CO-to-H$_2$ ratio, this lone cloud represents $\\sim10^{11}M_\\odot$ of cold and/or subthermally excited gas. In this paper, we present an analysis of nuclear CO(1-0) emission in \\apm\\ using VLA~\\footnote{The VLA is operated by the National Radio Astronomy Observatory, which is a facility of the National Science Foundation, operated under cooperative agreement by Associated Universities, Inc.} at high spatial resolution (0.3 arcsec). The CO appears as a partial ring of $\\sim$0.6 arcsec diameter. These data suggest a total revision in the gravitational lens model for this source, with the new model involving a `naked cusp', which naturally accounts for the observed odd-number of images. They also imply that the nuclear CO must be spatially extended on a scale of at least 400 pc, making this the second source in which gravitational lensing can be used as a `telescope' to explore sub-kpc scale structure of molecular gas in the AGN host galaxy. \\begin{figure*} \\centerline{ \\psfig{figure=MB659fig1.ps,angle=270,width=4.0in}} \\caption[]{The 23Ghz continuum (left) and CO(1-0) line emission (right) in \\apm. The contours are at -1 (dashed),0,1... $\\times \\sigma$. The image has been CLEANed and the ellipse in the upper corner of the left-hand panel represents the CLEAN beam. The quasar A image, as determined from NICMOS imaging, lies at 08 31 41.64, 52 45 17.5, offset from the radio emission (Ibata et al. 1999). This is probably due to astrometric errors in the VLA and HST reference frames.} \\label{fig1} \\end{figure*} ", + "conclusions": "This paper has presented resolved images of nuclear CO(1-0) emission in the gravitationally lensed BAL quasar \\apm. While the continuum emission is found to be well aligned with the optical quasar images, the CO(1-0) is more extended, with a broken ring-like appearance. Such a structure is consistent with the action of gravitational lensing, with the continuum emission occurring on the scale of the quasar core, while the CO(1-0) arises from a larger region and is differentially magnified. The three-image nature of \\apm\\ has posed a problem for lens modeling, as an extremely large, flat core is required to produce the central image. Such three image configurations are a nature consequence of gravitational lensing by a flattened potential which can produce naked cusps. Modeling of the CO(1-0) emission supports this hypothesis, although a deficit in constraints implies that the model is not unique. An immediate prediction of this model is that the lensing galaxy, whose position could be revealed by observing below the Lyman limit for this system ${\\rm (\\lta 4400\\AA)}$, hence removing the glare from the quasars, should be offset $\\sim0.5$arcsec from the quasar image, rather than lying behind the quasar images. Currently, our CO images of \\apm\\ are of limited signal-to-noise. However, with further integration a detailed map of the CO image can be made. As this region will be free from the effects of microlensing, and as its extended nature provides many more constraints (Kochanek, Keeton \\& McLoed 2001), such imaging has the potential to provide a more accurate model of the lensing in \\apm\\ than from the quasar images." + }, + "0112/astro-ph0112505_arXiv.txt": { + "abstract": "Radio Cherenkov radiation is arguably the most efficient mechanism for detecting showers from ultra-high energy particles of 1 PeV and above. Showers occuring in Antarctic ice should be detectable at distances up to 1 km. We report on electromagnetic shower development in ice using a GEANT Monte Carlo simulation. We have studied energy deposition by shower particles and determined shower parameters for several different media, finding agreement with published results where available. We also report on radio pulse emission from the charged particles in the shower, focusing on coherent emission at the Cherenkov angle. Previous work has focused on frequencies in the 100 MHz to 1 GHz range. Surprisingly, we find that the coherence regime extends up to tens of Ghz. This may have substantial impact on future radio-based neutrino detection experiments as well as any test beam experiment which seeks to measure coherent Cherenkov radiation from an electromagnetic shower. Our study is particularly important for the RICE experiment at the South Pole. ", + "introduction": "Ultra high energy (UHE) neutrinos can travel without scattering over large distances. These may prove to be useful cosmological and astrophysical probes. They also present themselves as candidate high energy particles with which we can test the Standard Model of electro-weak theory beyond the energy regime of current accelerators. In an UHE electron-neutrino charged current interaction, the neutrino gives most ($\\approx 80\\%$) of its energy to the secondary electron, which can then initiate an electromagnetic cascade or shower. It was predicted that an electromagnetic shower generated by a high energy primary could develop a charge excess which would emit Cherenkov radiation coherently \\cite{askaryan62}. For ultra high energy primaries the Cherenkov radiation would be coherent in the radio region of the spectrum; this long wavelength radiation might then be detected using radio antennas \\cite{markov86, provorov95, ralmc89}. Given the small predicted flux of ultra high energy neutrinos \\cite{Stecker96, halzen97, rachen99}, a suitable experiment to detect ultra high energy neutrinos using this method requires a large, dense (and radio-transparent) volume. Antarctic ice is suitable for this purpose. A detailed analysis of all aspects of such an experiment was done by Frichter, Ralston and McKay (FRM) \\cite{fmr96}, using a simulation developed by Zas, Halzen and Stanev (ZHS) \\cite{zhs92}. FRM concluded that a modest array of optimally designed antennas could detect many events per year. The sensitivity of radio detection peaks above 1 PeV, which compliments an optical array such as AMANDA \\cite{price96}. The Radio Ice Cherenkov Experiment (RICE) at the South pole \\cite{rice99} is designed as a prototype detector of ultra high energy neutrinos with energy $\\ge$PeV using this method. One basic requirement for such an experiment is a reliable Monte Carlo simulation of the shower development, Cherenkov radiation, detector, and data acquisition system. One can also test the idea of coherent Cherenkov emission at accelerator facilities by dumping a beam of photons or electrons into a dense target like sand or salt or any other suitable medium. Such tests have begun with experiments at Argonne and SLAC \\cite{saltzberg01}. A Monte Carlo simulation which can be easily adapted to such a test beam experiment, where Fresnel and possibly near zone radiation is important, and one that can include hadron showers conveniently, is clearly desirable. The ZHS simulation, designed for electromagnetic showers and Fraunhoffer (far zone) detection, has been a powerful tool. However an expansion and update with extensive testing, offering applications to test beam and neutrino astronomy, is currently needed. We have written a GEANT-based Monte Carlo simulation to study coherent Cherenkov emission in ice, salt, or a beam dump. GEANT 3.21 \\cite{geantman} is a well known and widely used simulation and detection Monte Carlo package in particle physics\\footnote{Differences between GEANT 3.21 and GEANT 4 are primarily at energies below the threshold for emission of Cherenkov radiation, and therefore do not affect the results presented here.}. It allows access to all details of the simulation such as controls of various processes, definition of target and detector media, and a complete history of all events simulated. GEANT can be used to simulate all dominant processes in 10 keV - 10 TeV energy range, although it has not been extensively verified for energies $>$100 GeV, where the extrapolation of well-known lower-energy electromagnetic cross-sections becomes large, and other effects (LPM, e.g. \\cite{landau53-1,landau53-2}) become significant. For electron energies above 10 GeV, GEANT uses screened Bethe-Heitler cross-section for {\\it bremsstrahlung} together with the Migdal corrections \\cite{messel70,migdal56}. The first Migdal correction is important for energy $\\ge$1 TeV, reducing the cross-section. The second correction reduces the differential cross-section for soft photon emission and is effective even at much lower energies, in the 100 MeV - 1 GeV range. The LPM effect in the context of UHE electromagnetic shower development and radio emission has been discussed in \\cite{misaki89, ralmc90, muniz97}. GEANT is used to simulate electromagnetic showers inside materials from which we extract detailed track information including particle type, coordinates, energy and interaction time. From this track information, we investigate shower properties like radiation length, Moliere radius, critical energy and energy deposition in the material. We also determine the resulting radio pulse using standard electrodynamic calculations from charged particles' tracks and by parametrizing the shower. Ultimately we will consider hadronic shower information and GEANT provides the flexibility to expand our analysis to this case. We note that physics results presented thus far by RICE have neglected the hadronic shower contribution and are, in this respect, conservative. The organization of this report is as follows. In section two, we discuss various aspects of electromagnetic showers and define quantities which characterize the shower. We present results on the shower structure from the Monte Carlo simulation in section three and compare them with established values in standard materials such as iron, lead and carbon. Our analysis includes the detailed breakdown of the radiation-generating charge imbalance into energy bins, the direct evaluation from energy considerations of Moliere radius and the determination from $\\dd E/\\dd x$ of radiation and energy deposition of the critical energy. We discuss shower-to-shower fluctuations by parametrizing the showers in section four. In section five, we review the theory behind coherent emission of an electromagnetic pulse from the shower \\cite{allan}. In section six, we calculate the electric pulse from the shower using track information from GEANT. We summarize our results in section seven, making a number of comparisons with ZHS, and discuss future work. ", + "conclusions": "We have analyzed 100 GeV - 1 TeV electromagnetic showers and the radio frequency radiation they produce in great detail. These studies are a necessary ingredient for designing an experiment for radio detection of UHE cosmic ray induced showers in radio- transparent media. In particular, experiments to detect radio emission from showers induced by high energy cosmic ray neutrinos interacting in the surface of the moon \\cite{gorham01} and in the South Polar ice-cap \\cite{rice01} are underway and have reported preliminary results. Coherent radio emission from electromagnetic showers has recently been demonstrated in the laboratory \\cite{saltzberg01}. The technique is gaining recognition as a powerful tool for particle detection. The thorough dissection and understanding of all the intricacies of the showers and their relationships to the final radio pulse produced is our goal achieved in this paper. Energy information for each stage of every track is readily available from the GEANT shower code. This information is essential for tracking the energy distribution in the shower and identifying the sources of radio emission. Among our results, is the {\\it direct} determination of the radiation length in ice from an exponential fit to bremsstrahlung radiation energy loss as a function of depth in Sec. 3.1. When this information is combined with the {\\it direct} extraction of the ionization loss, Sec. 3.3, we determined the {\\it critical energy} from the data shown in Fig. \\ref{fig:critenergy}. The value obtained is nicely consistent with that found by using the {\\it Moliere} radius extracted from the data for the radial energy flow, Sec. 3.2, in combination with the radiation length. The consistency between the critical energy values indicates that our application of the code and analysis of the data gives a correct physical picture of the interplay among the competing processes in the shower as it develops. Because the coherent radio emission of interest depends upon the charge excess in the shower, we need the energy profile of the contributions to the charge excess. This again requires the GEANT track-by-track energy information, and we show the total charge imbalance broken down into energy ranges in Fig. \\ref{fig:depfrac}. A large fraction, more than 50\\%, of the imbalance comes from the energy range below 5 MeV. Though the track lengths are small, the large number of particles leads to a significant contribution to the total track length of the shower. The total and projected track lengths for both the total and excess charge populations are determined and shown to be proportional to total shower energy in Sec. 3.5. The total track length is comfortably below a rough estimate of the upper bound for the total track length. Our detailed study of the longitudinal profiles in iron shows overall agreement but differences in detail among the GEANT, the EGS4 and the ZHS simulations in Fig. \\ref{fig:depiron}. The profiles in ice for 100 GeV, 500 GeV and 1 TeV electron and photon induced showers in ice are well described by a modified Greisen parametrization with critical energy value extracted from the simulation data, as described above, and two fit parameters. The confidence levels of the fits are typically 80\\% - 90\\%, as summarized in Tables 1 \\& 2. The GEANT profiles in ice are qualitatively similar to those from the ZHS code, but lie typically 25\\% - 35\\% lower. To determine whether small differences in cross section values used in the different simulations could account for this difference, we developed a 1-dimensional shower code with the full set of cross sections for the relevant processes. As we report in Sec. 3.7, the depth at maximum and the number of particles at maximum are rather insensitive to changes in the cross sections. We therefore believe that the differences in cross-sections are not the source of profile and total track length differences between our simulation and that of ZHS. Unfortunately, one important check that we have not been able to make is the GEANT vs. ZHS ionization energy loss ($\\dd E/\\dd x$). The ZHS code does not admit a readout of energy loss by shower particles. Therefore, we were unable to make from ZHS code plots similar to Figs. \\ref{fig:dedxdat} \\& \\ref{fig:dedxavg} we made with GEANT. If the ionization loss in ZHS code is much lower than in GEANT, it might account for the difference between showers produced by the two Monte Carlos. We develop the framework to calculate the electric field in the 100 MHz to multi GHz frequency range in Sec. 5. The track-by-track field calculation, applicable in the Fraunhoffer zone, and the calculation treating the shower as a continuous current density, applicable in the Fraunhoffer {\\it and} and Fresnel zones, complement each other. We present both for this reason. The direct, numerical calculation of the total field from the vector sum of the fields from individual tracks allows a detailed study of the dependence on total track length. Moreover, we elucidate the effects of random direction changes due to collisions and the pattern of phase relationships among the contributions from all tracks. Complementing these intensive numerical calculations, we present analytic work that employs the {\\it form factor} of the effective current density to calculate the field in both Fresnel and Fraunhoffer regimes. Given our realistic model for the current density, which we fit to the transverse charge distribution from the simulation, we show the remarkable result that the radiation from the shower is {\\it coherent over the whole shower} at the Cherenkov angle in the Fraunhoffer limit. Our study of the track-by-track phase coherence in Sec. 6 reveals that the phase associated with the initial time and position of a track (TP) and the phase associated with the ``diffractive emission'' (CP) from the track as a whole are uncorrelated. This result supports our model of the current, and consequently of the field, in Sec. 5. Figs. \\ref{fig:phase-c} \\& \\ref{fig:phase-40} show the strong phase peaking of the ``diffractive'' phase, called Cherenkov phase (CP) in the discussion, on and off the Cherenkov angle as defined by the shower axis. The phase associated with the initial coordinates of the track, called the translation phase (TP), is coherent at the Cherenkov angle but completely random away from that angle. We noticed some time ago \\cite{soebtalk00} that the frequency spectrum of the electric field at the Cherenkov angle, calculated directly from the track data and the Fraunhoffer zone formulas developed in Sec. 5., flattens out at frequencies above 2 GHz, as shown in Fig. 24. It is clear from the figure that the ZHS simulation shows the same behavior but they did not examine the high frequency region further \\cite{zhs92}. We employ the form factor, which correctly accounts for the transverse spread of the shower and its effect on the field to analyze the frequency spectrum. We evaluate its Fourier transform, $F(\\omega)$, two independent ways and find agreement. We then show that the observed spectrum of the field is faithfully represented up to 5 GHz for 100 GeV showers and up to 15 GHz for 500 GeV and 1 TeV showers. Up to 15 GHz, we now have a clear picture of the behavior of the frequency spectrum, which continues to show coherent behavior. The electric field in the Fraunhoffer zone as a function of the observation angle from the shower axis peaks at the Cherenkov angle. The width of the peak shrinks inversely with frequency and the height (field strength) rises linearly with shower energy. These features confirm the ZHS results and the off-Cherenkov angular dependence is also similar. The height of the peak at the Cherenkov angle for a given energy and frequency is larger in their case by about 25-35\\%. This is perhaps not a surprising difference between two independent shower simulations and field calculations. We have performed an extensive set of tests and cross-checks to validate our results. \\subsection{Conclusions} Our study quantifies coherent Cherenkov radiation from high energy showers and shows that coherence persists from 100 MHz to tens of GHz. The existence of Coherent Cherenkov radiation goes back to Askaryan, and has been studied for decades, while the persistence of coherent emission at the Cherenkov angle at multi-GHz frequencies is new. The persistence of coherence is established by our new, highly detailed study of the actual shower currents and corresponding phase distributions in CP and TP introduced in Sec. 6. The recognition that the TP is coherent over the whole length of the shower at the Cherenkov angle is a new insight into the connection between shower particles and the fields they produce. We have analyzed the electromagnetic shower characteristics in great detail, including the energy distributions and energy losses, so that a complete set of shower parameters can be extracted from the simulation data to confirm and cross check the consistency of our picture. A grand summary of pertinent parameters and comparisons with the ZHS and other sources where available is presented in Table 5. \\begin{center} \\begin{tabular}{lcc} \\multicolumn{3}{l}{Table 5: Comparisons between 100 GeV showers in ice from GEANT} \\\\ \\multicolumn{3}{l}{and from ZHS Monte Carlos}\\\\ \\hline \\hline Parameter & GEANT & ZHS \\\\ \\hline Total Energy Threshold (MeV) & 0.611 & 0.611 \\\\ Total absolute track length (meter) & 399 $\\pm$ 5 & 642 \\\\ Total projected ($e+p$) track length (meter) & 374 $\\pm$ 4 & 519 \\\\ Total projected ($e-p$) track length (meter) & 70 $\\pm$ 8 & 131 \\\\ Position of the shower max. (radiation length) & 6.5 & 7.0 \\\\ Number of particles ($e+p$) at shower max. & 111 $\\pm$ 7 & 155 $\\pm$ 10 \\\\ Excess electrons ($e-p$) at shower max. & 20 $\\pm$ 2 & 37 $\\pm$ 3 \\\\ Fractional charge excess at the shower max & $\\sim$ 18\\% & $\\sim$ 24\\% \\\\ Cherenkov peak at 1 GHz (Volts/MHz) & 7.5 $\\times$ 10$^{-9}$ & 1.0 $\\times$ 10$^{-8}$ \\\\ \\hline \\hline \\end{tabular} \\end{center} The simulation and field calculation developed are directly applicable to energies up to 1 TeV. Higher energy showers can be \"bootstrapped\" by evolving the shower particles into this regime with a corresponding multiplication in particle number and emitted field intensity. We intend to parametrize and extrapolate the data to the multi-PeV data needed for the UHE showers relevant to the RICE experiment. With the GEANT base, it will also be interesting to investigate similar shower and field calculation for the hadronic component of neutrino-induced showers. \\subsection*{Acknowledgement} Help and advice from Enrique Zas at various stages of this work were invaluable. Discussions and working sessions with Jaime Alvarez-Muniz were essential to our understanding of the workings of the ZHS code. We thank George Frichter and Florian Hardt for early diagnostic works using the ZHS code. Comments and questions from Frances Halzen on this topic over many years have been most useful. Thanks to Tim Bolton for helping with initial set up of the GEANT Monte Carlo code. This work is supported in part by the NSF, the DOE, the University of Kansas General Research Fund, the RESEARCH CORPORATION and the facilities of the Kansas Institute for Theoretical and Computational Science. \\newcommand{\\noopsort}[1]{} \\newcommand{\\printfirst}[2]{#1} \\newcommand{\\singleletter}[1]{#1} \\newcommand{\\switchargs}[2]{#2#1}" + }, + "0112/astro-ph0112443_arXiv.txt": { + "abstract": "An important and widely neglected aspect of the interaction between an accretion disc and a massive companion with a coplanar orbit is the vertical component of the tidal force. As shown by Lubow, the response of the disc to vertical forcing is resonant at certain radii, at which a localized torque is exerted, and from which a compressive wave (p mode) may be emitted. Although these vertical resonances are weaker than the corresponding Lindblad resonances, the $m=2$ inner vertical resonance in a binary star is typically located within the tidal truncation radius of a circumstellar disc. In this paper I develop a general theory of vertical resonances, allowing for non-linearity of the response, and dissipation by radiative damping and turbulent viscosity. The problem is reduced to a universal, non-linear ordinary differential equation with two real parameters. Solutions of the complex non-linear Airy equation are presented to illustrate the non-linear saturation of the resonance and the effects of dissipation. It is argued that the $m=2$ inner vertical resonance is unlikely to truncate the disc in cataclysmic variable stars, but contributes to angular momentum transport and produces a potentially observable non-axisymmetric structure. ", + "introduction": "\\subsection{Horizontal and vertical tidal forces} Accretion discs in binary stars and protoplanetary systems are subject to periodic tidal forcing from the orbiting companion objects. Similar effects can occur in planetary ring systems, where the planetary satellites provide the perturbing forces, and in galactic discs. Even in the simple case of a companion with a circular orbit that is coplanar with the disc, the effects of the tidal force are manifold. Consider a system of two spherical bodies of masses $M_1$ and $M_2$ in a bound orbit. Consider a third, test body orbiting about $M_1$. In a non-rotating coordinate system centred on $M_1$, the force per unit mass experienced by the test body is \\begin{equation} \\bff=-{{GM_1\\,\\br}\\over{|\\br|^3}}- {{GM_2(\\br-\\bd)}\\over{|\\br-\\bd|^3}}- {{GM_2\\,\\bd}\\over{|\\bd|^3}}, \\end{equation} where $\\bd(t)$ is the position vector of the companion $M_2$ with respect to $M_1$, and the third term is the fictitious force arising from the acceleration of the origin of the coordinate system. Inasmuch as the second and third terms perturb the Keplerian motion of the test body about $M_1$, they represent the tidal force in this system. Consider a thin accretion disc around $M_1$, and let $(R,\\phi,z)$ be cylindrical polar coordinates such that $z=0$ is the mid-plane of the disc. Suppose that the companion has a circular, coplanar orbit of radius $d$ and angular velocity $\\omega$. Then the components of the tidal acceleration are \\begin{equation} f_{{\\rm t}R}=-GM_2[R^2+d^2-2dR\\cos(\\phi-\\omega t)+z^2]^{-3/2} [R-d\\cos(\\phi-\\omega t)]-GM_2d^{-2}\\cos(\\phi-\\omega t), \\end{equation} \\begin{equation} f_{{\\rm t}\\phi}=-GM_2[R^2+d^2-2dR\\cos(\\phi-\\omega t)+z^2]^{-3/2} d\\sin(\\phi-\\omega t)+GM_2d^{-2}\\sin(\\phi-\\omega t), \\end{equation} \\begin{equation} f_{{\\rm t}z}=-GM_2[R^2+d^2-2dR\\cos(\\phi-\\omega t)+z^2]^{-3/2}z. \\end{equation} The horizontal components of the tidal acceleration are nearly independent of $z$ within the disc and cause a general non-axisymmetric distortion of its streamlines and surface density distribution (Papaloizou \\& Pringle 1977). Indeed, the streamlines of the disc correspond closely to a family of simple periodic orbits of the restricted three-body problem found by Paczy\\'nski (1977). The tidal distortion is enhanced in the neighbourhood of Lindblad resonances, radii at which the forcing frequency resonates with the natural epicyclic oscillations of the disc. Usually at a Lindblad resonance, a non-axisymmetric wave will be launched and will then propagate some way radially through the disc before dissipating and transferring its angular momentum to the disc (Goldreich \\& Tremaine 1979). The resonant torque exerted between the companion and the disc may be calculated from a standard formula that applies under quite general circumstances in linear theory, for example if the response is dominated by viscosity and no wave is emitted (Meyer-Vernet \\& Sicardy 1987), or if the vertical structure of the disc is taken fully into account (Lubow \\& Ogilvie 1998). Tidal torques usually limit the outer radius of a circumstellar disc (or the inner radius of a circumbinary disc). Except in the case of a very low-mass companion, the Lindblad resonances are so strong that they are excluded from the disc. In a binary star with a mass ratio $q=M_2/M_1$ of order unity, a circumstellar disc is tidally truncated well inside the innermost ($m=2$) Lindblad resonance, which in any case lies outside the Roche lobe (Paczy\\'nski 1977; Papaloizou \\& Pringle 1977). For smaller mass ratios typical of giant planets orbiting stars, the Lindblad resonances play a more direct role in truncating the disc (Lin \\& Papaloizou 1986). Another horizontal tidal effect is the eccentric interaction between the companion and the disc. At eccentric Lindblad resonances such as the 3:1 resonance, a local eccentric instability occurs, which may be able to sustain a large-scale eccentric distortion of the disc (Lubow 1991). The eccentric corotational resonances act to damp eccentricity, however. Studies of the {\\it vertical\\/} component of the tidal force have mostly been concerned with the case of a companion with an inclined orbit. In that case, the vertical tidal acceleration is nearly independent of $z$ within the disc. As a result, non-axisymmetric bending waves are launched at vertical resonances where the forcing frequency resonates with the natural vertical oscillations of the disc (Shu, Cuzzi \\& Lissauer 1983). In a Keplerian disc the Lindblad and vertical resonances coincide. For a companion with a {\\it coplanar\\/} orbit, however, the vertical tidal acceleration is proportional to $z$ and vanishes on the mid-plane. Other than those concerned with bending waves, most analytical and numerical studies have relied on a two-dimensional treatment of the disc, and the vertical component of the tidal force has been widely neglected. However, Lubow (1981) analysed the response of a vertically isothermal disc to vertical tidal forcing by a companion with a circular, coplanar orbit. The general response consists of a non-axisymmetric vertical expansion and contraction of the disc, characterized by a vertical velocity proportional to $z$. This `breathing' motion corresponds to a compressive mode of the disc, the $n=1$ p mode in the notation of Lubow \\& Pringle (1993), or the ${\\rm p}_1^{\\rm e}$ mode in the classification of Ogilvie (1998). At particular radii where the forcing frequency resonates with the natural frequency of this mode, a non-axisymmetric p-mode wave may be launched and a resonant torque exerted. These are also {\\it vertical resonances}, but are distinct from those associated with bending waves because a different mode is involved. Lubow's vertical resonances are further from the companion than the corresponding Lindblad resonances and, in a thin disc, the resonant torques are much weaker. However, Lubow (1981) pointed out that, in a binary star with a mass ratio of order unity, the $m=2$ inner vertical resonance typically lies within the standard tidal truncation radius, and the associated torque can compete with the viscous torque in the disc. Another consequence of the vertical tidal force is a local tilt instability, closely related to the eccentric instability (Lubow 1992). The tilt instability occurs at inclination resonances that are essentially coincident with the eccentric Lindblad resonances, but the tilt instability is much weaker. \\subsection{Aims of the paper} In this paper I develop a non-linear theory of Lubow's vertical resonances. I generalize the analysis of Lubow (1981) to non-isothermal discs allowing for non-linearity, viscosity and radiative damping. By reducing the problem to a universal differential equation and studying its solutions, I aim to demonstrate the effects of non-linear saturation and dissipation on the resonant response of differentially rotating discs to periodic forcing in general. In a related paper (Ogilvie 2001, hereafter Paper~I) I have presented an analysis of the non-linear tidal distortion of a thin, three-dimensional accretion disc by a binary companion on a circular orbit. I showed that the resulting distortion can plausibly explain the two-armed non-axisymmetric features seen in the Doppler tomograms of IP~Peg and other dwarf novae in outburst (e.g. Steeghs 2001). Although the $m=2$ inner vertical resonance plays an important role in regulating the amplitude and phase of the tidal distortion, the method of analysis in Paper~I does not allow for the emission of a wave from the vertical resonance. One of the aims of the present paper is to justify that limitation by showing that, under typical conditions, vertical resonances are broadened and damped by non-linearity and dissipation to the extent that the emission of a wave may reasonably be neglected. Before embarking on the detailed analysis, it may be helpful to look ahead to the equation to be derived, \\begin{equation} -{{{\\rm d}^2y}\\over{{\\rm d}x^2}}+xy+|y|^2y+{\\rm i}by=a. \\end{equation} In this equation, $x$ represents the radial distance from the resonant orbit and $y$ the complex amplitude of the response, both in dimensionless terms. The linear terms $-y''+xy$ represent the ability of the disc to support freely propagating waves in the region $x<0$. In the present case these waves correspond to the compressive ${\\rm p}_1^{\\rm e}$ mode. The point $x=0$ is the location of the resonance, which is also the turning point of free waves. The constant term $a$ on the right-hand side represents the tidal forcing in the neighbourhood of the resonance. The inhomogeneous Airy equation $-y''+xy=a$ is familiar in studies of resonant wave excitation in differentially rotating discs and features in the analysis of Lubow (1981). The new, non-linear term $|y|^2y$ derives from couplings between the resonant mode and non-resonant modes. It leads to new effects such as the non-linear broadening and saturation of the resonance. The correct derivation of this term accounts for much of the complexity of the analysis in the present paper. The linear term ${\\rm i}by$ represents the dissipation of vertical motions by shear and bulk viscosity and by radiative damping. It leads to attenuation of the waves and broadening of the resonance. The remainder of this paper is organized as follows. In Section~2 I recall the coordinate system and the basic equations used in Paper~I. In Section~3 I expand the equations in the neighbourhood of the vertical resonance and solve them in a systematic manner to derive the complex non-linear Airy equation. In Section~4 I discuss the properties of the equation and present numerical solutions for a range of parameter values. The astrophysical consequences of the analysis are explored in Section~5. ", + "conclusions": "In this paper I have developed a general theory of vertical resonances, first analysed by Lubow (1981), which are an important aspect of the interaction between an accretion disc and a massive companion with a coplanar orbit. When Lubow's analysis is generalized to allow for non-linearity of the response, and dissipation by radiative damping and turbulent viscosity, the problem reduces to a universal, non-linear ordinary differential equation with two real parameters. Numerical solutions of a time-dependent version of the complex non-linear Airy equation describe the process by which a steady and stable pattern of non-axisymmetric distortion is established in an initially undisturbed disc. For small values of the dissipation parameter $b$, a p-mode wave is launched at the resonance and propagates radially through the disc, carrying angular momentum. For larger values of $b$ the wave is attenuated or not launched at all, and the resonant response is broadened into a smooth feature. The effect of increasing the amplitude parameter $a$ is to broaden the resonance and ultimately to saturate it. The total tidal torque exerted at the resonance is independent of $b$ in the linear regime $a\\ll1$, but can be either greater or less than the linear prediction when non-linearity is important. If the magnitude of the resonant tidal torque exceeds that of the local viscous torque, the disc may be truncated at the resonance (Lubow 1981). The viscous torque in the absence of tidal distortions is \\begin{equation} \\int\\!\\!\\!\\int\\lambda^3\\mu{{{\\rm d}\\Omega}\\over{{\\rm d}\\lambda}} \\,{\\rm d}\\phi\\,{\\rm d}z= -3\\pi\\alpha\\lambda^2\\tilde\\Omega^2\\int\\rho z^2\\,{\\rm d}z, \\end{equation} and a comparison with equation (\\ref{torque}) shows that truncation occurs if $\\alpha<\\alpha_{\\rm trunc}$, where \\begin{equation} \\alpha_{\\rm trunc}={{f_Tm\\Psi^2}\\over{3\\sD\\tilde\\Omega^2}}. \\end{equation} However, if $\\alpha$ is so small that the launched wave does not damp significantly in the vicinity of the resonance, the tidal torque is not transmitted to the disc locally and the process of truncation cannot occur in a straightforward manner. In the derivation of the complex non-linear Airy equation it was assumed that $\\alpha=O(\\epsilon^{2/3})$ while $\\Psi=O(\\epsilon)$. Therefore the process of tidal truncation is not described by that equation, but occurs in a different parameter regime. Some illustrative values of $\\alpha_{\\rm trunc}$ for the $m=2$ inner vertical resonance are given in Table~1, where it is assumed that $f_T=1$. It appears that truncation by this resonance is unlikely in cataclysmic variable discs, even in quiescence, and especially if $\\gamma\\approx5/3$. Although Lubow (1981) argued that truncation is likely, his estimates were based on the case $\\gamma=1$, for which the resonant radius is maximal. This results in a much larger torque than is obtained when $\\gamma=5/3$, because of the proximity to the companion. However, for the same reason, the resonance is less likely to lie inside the disc in the case $\\gamma=1$. \\begin{table*} \\centering \\caption{Critical value of the viscosity parameter below which the disc may be truncated at the $m=2$ inner vertical resonance.} \\begin{tabular}{@{}lrr@{}} $q$&$\\gamma$&$\\alpha_{\\rm trunc}$\\\\ \\hline 1&1&0.07489\\\\ &6/5&0.01694\\\\ &5/3&0.00037\\\\ 0.5&1&0.03533\\\\ &6/5&0.00796\\\\ &5/3&0.00018\\\\ 0.2&1&0.00935\\\\ &6/5&0.00210\\\\ &5/3&0.00005 \\end{tabular} \\end{table*} Much more probable is that a non-destructive, two-armed structure is formed, as described by the complex non-linear Airy equation. In Paper~I I have argued that this effect, in the limit in which the resonant response is broadened and damped by non-linearity and dissipation, could explain the non-axisymmetric structures seen in Doppler tomograms of dwarf novae in outburst. Indeed, estimates of the parameters for a system such as IP~Peg (Section~\\ref{illustrative}) suggest that the resonance is strongly damped and mildly non-linear. It is likely that the techniques used in this paper will be useful in other circumstances where spatial resonances occur in wave-bearing media subject to periodic forcing. The equation derived has a concise form and its terms have a clear physical interpretation, suggesting that it may be generic to some extent. In any case it provides a simple mathematical model in which the effects of non-linearity and dissipation on the resonant launching of waves can be studied." + }, + "0112/astro-ph0112169_arXiv.txt": { + "abstract": "We explore the metallicity distribution function (MDF) of red giant stars in $\\omega$ Centauri from a catalogue of Washington $M$, $T_2$ and $DDO51$ photometry covering over 1.1 deg$^2$ outside the cluster core. Using updated calibrations of giant branch isometallicity loci in this filter system, photometric metallicities, guided by previously published spectroscopic abundances, are derived. Several methods are employed to correct the MDF for contamination by Galactic stars, including: (1) use of the surface gravity sensitivity of the ($M-DDO51$) color index to eliminate foreground dwarf stars, (2) radial velocities, and (3) membership probabilities from proper motions. The contamination-corrected MDF for $\\omega$ Cen shows a range of enrichment levels spanning nearly 2 dex in [Fe/H], and with peaks at [Fe/H]=$-1.6$, $-1.2$, and $-0.9$. ", + "introduction": "The large abundance spread seen in the red giant branch (RGB) of $\\omega$ Cen has long been recognized as one of the unique features of this peculiar Milky Way globular cluster. Recent photometric analyses of the $\\omega$ Cen RGB (e.g., Lee et al.\\ 1999; Pancino et al.\\ 2000, PFBPZ hereafter; Majewski et al.\\ 2000a, M00a hereafter) indicate a metallicity distribution function (MDF) stretching from [Fe/H]$\\sim-2.0$ to perhaps as high as [Fe/H]=$-0.4$. This spread, together with clear evidence for a 2-4 Gyr age spread (Hughes \\& Wallerstein 2000) as well as other unusual characteristics relating to its large mass, elongated shape, and internal and external dynamics (see summary in M00a), suggests that $\\omega$ Cen may represent an important transitional link between globular clusters and dwarf galaxies. Here we revisit the M00a analysis of the $\\omega$ Cen MDF with the addition of new membership data for correcting sample contamination and an improved photometric metallicity calibration. ", + "conclusions": "It is well known that $\\omega$ Cen does not conform with most globular clusters in a variety of ways (see M00a): It is the most massive cluster, it shows substantial rotation and flattening, and, of course, it has a large metallicity spread. Several theories about the origin of $\\omega$ Cen have been proposed, including that it is a rare cluster that (for some reason) encountered substantial self-enrichment, that it is the product of the merger of two stellar systems, that it is the remains of a disrupted dwarf spheroidal, and even that it derived from some amalgam of these possibilities. That $\\omega$ Cen seems to have at least {\\it three} primary enrichment peaks and an overall [Fe/H] spread from at least $-$0.4 to $-$2.0 dex, coupled with claims for an age spread of up to 4 Gyr in the cluster's main sequence turn-off (Hughes \\& Wallerstein 2000), makes a simple two cluster merger hypothesis unlikely (see also Norris et al. 1997). Confronted by the difficulties of multiple metallicity populations and motivated by the relative spatial distributions of these populations, PFBPZ propose a more complicated scenario -- the merger of two systems with at least one of the systems having undergone self-enrichment and sinking into the center of $\\omega$ Cen. For the merged, self-enriched entity, which is intended to account for the two intermediate as well as the most metal rich populations, PFBPZ propose a giant molecular cloud or a gas-rich protocluster. However, a number of aspects of $\\omega$ Cen lead one to suspect its closer association with dwarf galaxies. For example, the ``peaky\" MDF of $\\omega$ Cen bears great resemblance to the burst-like, multiple populations seen in dwarf spheroidal (dSph) galaxies (Grebel 1997). Interestingly, the Sagittarius (Sgr) dwarf galaxy shows a similarly large (and punctuated) spread in [Fe/H] to $\\omega$ Cen (Layden \\& Sarajedini 2000). For a variety of reasons, including the similarity in MDFs as well as the fact that the mass of $\\omega$ Cen is comparable to that of the globular M54, which appears to be the core of Sgr, it has been proposed (e.g., Lee et al.\\ 1999, M00a) that $\\omega$ Cen may be the remnant nucleus of a tidally disrupted dwarf galaxy analogous to the Sgr system. As pointed out by Shetrone et al. (2001), for this model of $\\omega$ Cen formation to work, the cluster would have to be a daughter product of a large dwarf galaxy like Sgr, since the heavy-element abundance patterns of smaller, dSph systems like Ursa Minor, Draco and Sextans differ from that of $\\omega$ Cen, which shows a large enhancement of s to r-process elements with increasing metallicity (Smith et al. 2000). On the other hand, the apparent greater concentration of more metal rich stars observed in $\\omega$ Cen by PFBPZ mimics a trend seen in dwarf galaxies both great (like Fornax -- Grebel \\& Stetson 1998) and small (like Sculptor -- e.g., Majewski et al.\\ 1999). Apart from the actual difficulty of two clusters merging, which requires relative velocities of $<\\sim 1$ km s$^{-1}$ (Thurl \\& Johnston, this proceedings), the merger hypothesis suffers from at least one other unlikelihood: If $\\omega$ Cen were the result of the merger of two cluster-like systems, the parent clusters would {\\it each} have to have been among the largest clusters in the Galaxy, and even if only the metal poor part of $\\omega$ Cen began its life as a traditional cluster, it too would be at the extreme end of the Galactic cluster mass scale. Somehow it is easier to accept that the peculiar properties of $\\omega$ Cen are the {\\it result} of its large mass, rather than that its large mass and other peculiar properties were accumulated as the result of a series of unlikely occurences. Indeed, the present orbit of $\\omega$ Cen (i.e., barreling retrograde within and through the Galactic plane -- Dinescu et al. 1999) is one that undoubtedly subjects it to substantial tidal stripping. Therefore, not only was $\\omega$ Cen almost certainly larger and even more like a dwarf galaxy in the past, but there is every expectation that it has led a battered life much like its Sgr counterpart. Evidence for tidal tails extending from $\\omega$ Cen have been reported by Leon et al. (2000). We have attempted to present a more accurate representation of the MDF for $\\omega$ Cen. However, as pointed out by Majewski et al. (2001), if a system has endured substantial mass loss over its lifetime, one must be wary of interpreting the presently observed MDF to represent the true enrichment history of that stellar system. Older (and more extended) populations will have had more time to have been stripped, and especially in the case of $\\omega$ Cen, whose planar orbit has almost certainly evolved considerably, that mass loss rate may have been highly variable over the enrichment timescale. We thank support from the National Science Foundation, The David and Lucile Packard Foundation, Research Corporation and Carnegie Observatories." + }, + "0112/astro-ph0112396_arXiv.txt": { + "abstract": "We develop the maximum-entropy weak shear mass reconstruction method presented in earlier papers by taking each background galaxy image shape as an independent estimator of the reduced shear field and incorporating an intrinsic smoothness into the reconstruction. The characteristic length scale of this smoothing is determined by Bayesian methods. Within this algorithm the uncertainties due to both the intrinsic distribution of galaxy shapes and galaxy shape estimation are carried through to the final mass reconstruction, and the mass within arbitrarily shaped apertures can be calculated with corresponding uncertainties. We apply this method to two clusters taken from $n$-body simulations using mock observations corresponding to Keck LRIS and mosaiced HST WFPC2 fields. We demonstrate that the Bayesian choice of smoothing length is sensible and that masses within apertures (including one on a filamentary structure) are reliable, provided the field of view is not too small. We apply the method to data taken on the cluster MS1054-03 using the Keck LRIS (Clowe et al. 2000) and HST (Hoekstra et al. 2000), finding results in agreement with this previous work; we also present reconstructions with optimal smoothing lengths, and mass estimates which do not rely on any assumptions of circular symmetry. The code used in this work (LensEnt2) is available from the web. ", + "introduction": "\\label{intro} Weak lensing studies of clusters of galaxies are an important complement to X-ray, Sunyaev-Zel'dovich effect and optical observations, allowing the projected distribution of mass to be investigated without any dynamical assumptions. The reconstruction of cluster mass distributions from weak gravitational lensing data is now well established; it has been shown that the projected density distribution can be recovered from magnification data, in the form of background galaxy number densities~\\cite{BTP95,DT98}, or from shear data, the net statistical distortion of the images of background galaxies~\\cite{TVW90,KS93,SS95,SK96}. We focus here on shear data, primarily because of its greater abundance; the likelihood function for shear data is also better understood~(Section~\\ref{method}). Schneider, King and Erben~\\shortcite{SKE00} discuss the use of the two types of weak gravitational lensing data. Reconstruction methods using shear data fall into two classes: direct and iterative inverse methods. The direct methods are based on the pioneering work of Kaiser and Squires~(1993, KS93); many improvements have since been made to the original algorithm~\\cite{SS95,Kai95,Bar95,SK96}. In all these methods the galaxy shape data have to be smoothed before their input to the algorithm; the smoothing length is a parameter that is left undetermined. The class of iterative methods aims to find the mass or projected gravitational potential map that best fits the data~\\cite{SK96,BNSS,SSB98}. These methods are well suited to irregularly shaped observations, since they do not suffer from edge effects in the same way as the direct methods; however, they need to be regularised in some way to prevent over-fitting the data, and it remains unclear how best to determine the resolution of either the data bins or the reconstruction grid. In two earlier papers~(Bridle~\\ea 1998, \\ppi; Bridle~\\ea 2001, \\ppii) we presented a maximum-entropy inverse method for reconstructing the mass distribution in clusters using shear and/or magnification data. In this paper we extend our method to give a fuller Bayesian analysis. As noted by other authors~\\cite{SSB98}, it would be desirable to work with each background galaxy shape individually, rather than binning or smoothing the data. This issue, together with the problem of the angular resolution of the reconstruction, is addressed by our extended algorithm. We apply our improved method to both realistic synthetic data, and previously published data for the high-redshift cluster MS1054-03. As with any Bayesian analysis, the aim is to derive and interpret the full posterior probability distribution of the quantity being inferred (in this case the mass distribution and any associated parameters). This approach will provide us not just with a mapping procedure, but also valuable insight into the quality of the data itself. The method is reviewed and further developed in Section~\\ref{method}, and is applied to simulated data in Section~\\ref{simobs}. Section~\\ref{realdata} contains the results of our method applied to the well-documented cluster MS1054-03, and gives a brief comparison with the previously published work. Our conclusions are presented in Section~\\ref{conc}. ", + "conclusions": "\\label{conc} We have developed a Bayesian analysis, based on the maximum-entropy method of Bridle~\\ea(1998, 2001), for inferring the distribution of mass in clusters of galaxies from weak lensing shear data. We treat each background galaxy image as a noisy estimator of the reduced shear field of the cluster, retaining all the information about both signal and noise and so allowing the for high angular resolution. Use of an `intrinsic correlation function' in the maximum-entropy formalism provides a way of incorporating our prior expectation of clusters as smooth, extended objects, and effectively replaces the data smoothing required by direct reconstruction methods. In contrast to the these methods, the lensing signal is not diluted by this process. Moreover, analysis of the posterior probability distribution of the ICF width~$w$, obtained by numerically evaluating the Bayesian evidence, provides an objective way of discriminating between smoothing scales. The map at the peak of this probability distribution was found not to contain any significant spurious peaks, and can be interpreted as the safest conclusion to draw from the data. The higher resolution maps, although representing an overfit to the data, do contain limited useful information particularly with respect to substructure in the cluster; the fact that their angular resolution is less favoured by the data quality gives a useful indication of the believability of these features. Simple mass estimates extracted directly from the mass maps preferred by the evidence were found to be unbiased and accurate to within the estimated errors over a fairly wide range of angular scales; these uncertainties agreed very well with the standard deviation in the mass estimates from 100 different realisations of the background galaxy population. Noisier observations over a smaller field of view were found to give mass estimates with a slight bias towards overestimation, an effect understood in terms of the additivity of the reconstructed distribution and the mass sheet degeneracy. Inspection of the variety of structure in reconstructions from these different noise realisations justified the cautious interpretation of the high resolution maps, indicating that this analysis provides a useful way of understanding the noise properties of the data. We have applied our method to two galaxy shape datasets for the high redshift cluster MS1054-03, one derived from a ground based observation and the other from an HST mosaic. In both cases the features found in previously published maps, obtained by both direct and inverse methods, are reproduced, but with the added desirable features of being positivity-constrained and quantitatively useful. Simple mass estimates extracted directly from the reconstruction agree well with values found by aperture mass densitometry, as do the errors estimated by both methods. The principal function of parameter-free mass maps as produced by this method is to provide the equivalent of a mass telescope, allowing images to be generated to aid further, more quantitative analysis. Parameterised fitting to shear data with physical motivation has been performed elsewhere~\\cite{SKE00,KS01}; we believe that the information provided by our variable angular resolution analysis is a helpful guide to this process, whilst also providing reasonably accurate `mass photometry' at the same time. The use of the Bayesian evidence in such fitting will be addressed in future papers. \\subsection*" + }, + "0112/astro-ph0112382_arXiv.txt": { + "abstract": "As is known, resonant scattering can distort the surface-brightness profiles of clusters of galaxies in X-ray lines. We demonstrate that the scattered line emission should be polarized and possibly detectable with near-future X-ray polarimeters. Spectrally-resolved mapping of a galaxy cluster in polarized X-rays could provide valuable independent information on the physical conditions, in particular element abundances and the characteristic velocity of small-scale turbulent motions, in the intracluster gas. The expected degree of polarization is of the order of 10\\% for richest regular clusters (e.g. Coma) and clusters whose X-ray emission is dominated by a central cooling flow (e.g. Perseus or M87/Virgo). ", + "introduction": "Although the intergalactic gas in clusters of galaxies is optically thin to Thomson scattering ($\\taut\\lesssim 0.01$) for the continuum photons, it can be moderately thick ($\\tau\\sim 1$) in the resonance X-ray lines of highly ionized atoms of heavy elements. This causes the radiative transfer in such lines to be important and gives rise to three major observable effects as previously discussed in the literature. First, the surface-brightness profile of a cluster in a resonance line, calculated in the optically-thin limit, should be distorted due to diffusion of photons from the dense core into the outer regions of the cluster (Gilfanov, Sunyaev \\& Churazov 1987a; Shigeyama 1998). This must be taken into account when attempting to determine element abundances from X-ray spectroscopic observations of galaxy clusters or haloes of elliptical galaxies. This effect has been already considered in application to ASCA, Beppo-SAX and XMM-Newton data \\cite{akietal97,tawetal97,moletal98,bohretal01,matetal01}. Resonant scattering also determines the spectral profile (which becomes saddle-shaped at small projected radii if $\\tau\\gtrsim 1$) of the emergent line \\cite{giletal87a}. With next generation X-ray spectrometers, combining high sensitivity with high energy resolution ($E/\\Delta E\\sim 10^3$--$10^4$), it should be possible to measure the profiles of individual lines and thus obtain a good deal of new information on the physical conditions in the intergalactic gas. Finally, resonant X-ray lines from the intracluster gas should be observable as absorption lines in the spectrum of a bright X-ray emitting background AGN. The strongest absorption lines can have equivalent widths of several eV (Shapiro \\& Bahcall 1980; Basko, Komberg \\& Moskalenko 1981; David 2000). By comparing the equivalent width of such a line and the flux from the same line seen in emission, the distance to the cluster can be directly determined, independent of the standard extragalactic distance scale (Krolik \\& Raymond 1988; Sarazin 1989). However, there should be yet another effect related to the problem under consideration. Namely, scattering in certain (as discussed below) resonance X-ray lines causes the scattered emission to be polarized. The recent advent of new technologies \\cite{cosetal01} that promise a drastically increased sensitivity of near-future X-ray polarimeters has motivated us to consider this effect in detail. ", + "conclusions": "\\begin{itemize} \\item We have shown that the expected degree of polarization for the brightest X-ray lines from clusters of galaxies is of the order of 10\\%. \\item The degree of polarization is a function of projected distance from the cluster center, and the polarization vector is perpendicular to the projected radius vector. This implies in particular that the integrated X-ray flux from the whole cluster will be unpolarized (provided the cluster if spherically symmetric). \\item It has also been demonstrated that for regular clusters, whose gas density distribution obeys the beta-law (\\ref{dens_prof}), the largest signal-to-noise ratio in measuring the polarization signal is attainable in the range of projected radii between $0.5$ and $10$ core-radii. The degree of polarization is zero in the direction of the cluster center, while the photon flux from the outer parts is small making it difficult to measure the polarization. \\item Rich regular clusters as well as clusters with a dominant cooling flow are the best targets for future polarimetric observations. \\end{itemize} Why is it important to observe the polarization effect? First of all, this effect is unique in the sense that there is no other obvious mechanism that could lead to polarization in X-ray lines. Thus, a detection of polarized diffuse emission in an X-ray line would be a solid proof that resonant scattering is indeed in operation. As mentioned before, resonant scattering affects the ``apparent'' (inferred under the assumption of optically thin emission) element abundance radial profiles of clusters \\cite{giletal87a,shigeyama98}. As a consequence, in order to determine the true abundance profiles from spectroscopic measurements, one should be able to accurately subtract the distortions caused by resonant scattering, which can well be larger in magnitude than the true underlying radial trend. This is a difficult task, because the surface brightness is an integral characteristic of a given line of sight which passes through regions with different densities and temperatures. One way to proceed is to compare the surface brightness profiles of at least two spectral lines corresponding to a given element (e.g. $K\\alpha$ and $K\\beta$ lines of iron, see Akimoto et al. 1997; Molendi 1998). Polarimetric measurements could significantly simplify such an analysis, for measuring of the polarized flux in a single line already allows one to estimate the contribution of resonant scattering to apparent abundance profiles. Apart from element abundances, the resonant scattering effects in principle can provide information on the velocity field within the cluster, because the optical depth in a given line is a function of the characteristic velocity of the corresponding ions. Lines of iron are particularly sensitive to the velocity of turbulent motions. Finally, the polarization effect can also be used for cosmological purposes. Indeed, since the surface brightness of resonantly scattered emission and of its polarized component is roughly proportional to the optical depth along a given line of sight while the surface brightness of continuum thermal emission is proportional to the integral of the density squared along the line of sight, it is possible to determine the distance to the cluster and consequently the Hubble constant analogously as is done using resonance X-ray absorption lines \\cite{kroray88,sarazin89} or the cosmic microwave background (the Sunayev-Zel'dovich effect). There are obvious requirements to a future X-ray polarimeter. First, it should have the sufficient angular resolution to probe individual parts of a cluster (we recall that the polarization vanishes when integrated over the whole cluster). Second, the CCD-type energy resolution is desirable in order to avoid significant contamination of the polarized emission of a resonance line by the unpolarized emission of neighbouring lines and the continuum. Finally, a large effective area of the telescope is needed, larger than is required for the detection of unpolarized resonantly scattered emission, since the expected degree of polarization is $\\sim 10$\\%." + }, + "0112/astro-ph0112457_arXiv.txt": { + "abstract": "\\vspace{\\baselineskip} We investigate the impact of peculiar velocity effects due to the motion of the solar system relative to the microwave background (CMB) on high resolution CMB experiments. It is well known that on the largest angular scales the combined effects of Doppler shifts and aberration are important; the lowest Legendre multipoles of total intensity receive power from the large CMB monopole in transforming from the CMB frame. On small angular scales aberration dominates and is shown here to lead to significant distortions of the total intensity and polarization multipoles in transforming from the rest frame of the CMB to the frame of the solar system. We provide convenient analytic results for the distortions as series expansions in the relative velocity of the two frames, but at the highest resolutions a numerical quadrature is required. Although many of the high resolution multipoles themselves are severely distorted by the frame transformations, we show that their statistical properties distort by only an insignificant amount. Therefore, cosmological parameter estimation is insensitive to the transformation from the CMB frame (where theoretical predictions are calculated) to the rest frame of the experiment. ", + "introduction": "\\label{sec:intro} The impressive advances being made in sensitivity and resolution of microwave background (CMB) experiments demand that careful attention be paid to potential systematic effects in the analysis pipeline. Such effects can arise from imperfect modelling of the instrument, e.g.\\ approximations in modelling the beam~\\cite{challinor00,wu01,souradeep01,fosalba01}, or incomplete knowledge of the pointing, but also from more fundamental effects such as inaccurate separation of foregrounds (see e.g.\\ Refs.~\\cite{bouchet99,hobson99} for reviews). In this paper we consider errors that may arise due to neglect of the peculiar motion of the experiment relative to the CMB rest frame (that frame in which the CMB dipole vanishes). For short duration experiments (e.g. balloon flights such as MAXIMA~\\cite{maxima} and BOOMERANG~\\cite{boomerang}) the relative velocity is constant over the timescale of the experiment, but for experiments conducted over a few months or longer, and particularly for satellite surveys~\\cite{cobe,map,planck}, the variation in the relative velocity adds additional complications. In principle, the modulation of the aberration arising from any variation in the relative velocity must be accounted for with a more refined pointing model for the experiment~\\cite{vanleeuwen02,challinor02} when making a map. For a relative speed of $\\beta c$ (where $c$ is the speed of light and $\\beta\\sim 1.23\\times 10^{-3}$ for the solar-system barycenter relative to the CMB frame), the r.m.s.\\ photon Doppler shifts and deflection angles are $\\beta/\\sqrt{3}$ and $\\sqrt{2/3}\\beta$ respectively. Despite these small values, significant distortions of the spherical multipoles of the total intensity and polarization fields do arise. A well known example is provided by the CMB dipole seen on earth, which, given the observed spectrum, arises from the transformation of the monopole in the CMB frame. More generally, on the largest angular scales the combined effects of Doppler shifts and aberration couple the total intensity monopole and dipole into the $l$th multipoles at the level $O(\\beta^l)$ and $O(\\beta^{l-1})$ respectively. Given the size of the non-cosmological monopole, annual modulation of the dipole by the variation in the relative velocity of the earth in the CMB frame must be considered in long duration experiments. In this paper we concentrate on the effects of peculiar velocities on small angular scale features in the microwave sky. On such scales, aberration dominates the distortions and becomes particularly acute when the angular scales of interest, $O(1/l)$, drop below the r.m.s.\\ deflection angle, i.e.\\ $l \\agt 800$ for the transformation from the CMB frame to that of the solar system. We provide simple analytic results for these distortions to the total intensity and polarization fields as power series in the relative velocity $\\beta$. The power series converge rather slowly at the highest multipoles for most values of the azimuthal index $m$ [the leading-order corrections go like $O(l\\beta)$] but the distortions can still easily be found semi-analytically with a one-dimensional quadrature. If the transformations of the multipoles carried through to their statistical properties, theoretical power spectra computed in linear theory (e.g.\\ with standard Boltzmann codes~\\cite{seljak96,LCL99}) would not accurately describe the statistics of the high resolution multipoles observed on earth. (The theoretical power spectra would still be accurate in the CMB frame.) It is straightforward to calculate the statistical correlations of the multipoles observed on earth. Fortunately, as we show here, the statistical corrections due to peculiar velocity effects turn out to be negligible despite the large corrections to the individual multipoles. It follows that for the purposes of high resolution power spectrum and parameter estimation, the transformation from the CMB frame can be neglected. This paper is arranged as follows. In Sec.~\\ref{sec:intensity} we describe the transformation laws for the total intensity multipoles in specific intensity and frequency-integrated forms. Convenient series expansions in $\\beta$ of the transformations are provided, and their properties under rotations of the reference frames are described. The statistical poperties of the transformed multipoles are investigated by constructing rotationally-invariant power spectrum estimators and full correlation matrices. In Sec.~\\ref{sec:polarization} we discuss the geometry of the frame transformations for linear polarization, and present power series expansions for the transformations of the multipoles. The behaviour under rotations and parity are also outlined. Power spectra estimators and correlation matrices are constructed, and cross correlations with the total intensity are considered. Some implications of our results for survey missions are discussed in Sec.~\\ref{sec:discussion}, which is followed by our conclusions in Sec.~\\ref{sec:conclusion}. An appendix provides details of the evaluation of the multipole transformations as power series in $\\beta$. We use units with $c=1$. ", + "conclusions": "\\label{sec:conclusion} We have shown that for total intensity the effect of the frame transformation from the CMB frame to that of the solar system produces large distortions in certain multipoles at high $l$. These effects arise principally from aberration rather than Doppler shifts. The linear polarization multipoles are similarly distorted at high $l$, but with the additional complication that there is some transfer of power between $E$ and $B$ polarization. This transfer is suppressed at large $l$, and receives comparable contributions from aberration and Doppler shifts on all scales. Although the power in $B$ polarization is expected to be much smaller than that in $E$ in the absence of foregrounds, the $B$ polarization generated from $E$ is well below the primordial level even if gravity waves contribute only one percent of the large-angle temperature anisotropies. If the gravity wave background is much below this level, weak gravitational lensing will dominate the primordial signal on all scales. This lensing signal is expected to be an order of magnitude larger than the $B$ polarization generated from the frame transformation on large scales. Despite significant $O(\\beta l)$ distortions of certain multipoles at large $l$, peculiar velocity effects are suppressed in power spectrum estimators and the covariance matrices for the CMB signals. The effect of the frame transformation on the mean of the simplest power spectrum estimator is to convolve the spectrum in the CMB frame (which we can compute reliably with linear perturbation theory) with a narrow kernel that integrates to unity. For smooth spectra there is negligible bias introduced by such a convolution. For linear polarization, the bias of e.g.\\ the $B$-polarization power spectrum by $E$ is suppressed at large $l$, and is expected to be negligible on all scales. We also showed that the frame transformation has only a negligible effect [$O(\\beta)$ as opposed to $O(\\beta l)$] on the signal covariance matrices for smooth underlying power spectra. The leading order effect is a coupling to the adjacent $l$ values, $l\\pm 1$. For linear polarization additional correlations are induced between $B$ and $E$ polarization, and $B$ and total intensity $I$, since the frame transformation does not preserve parity invariance, but their level is negligible. If the CMB fluctuations are Gaussian in the CMB frame, the multipoles will remain Gaussian distributed in any other frame since the transformation is linear in the signal. The transformation does break rotational and parity invariance, however, and so the aberration effects described here may be important when searching for weak lensing effects in the microwave background (using small patches of the sky over a coherence area of the weak shear), or the effects of non-trivial topologies." + }, + "0112/astro-ph0112511_arXiv.txt": { + "abstract": "{Compact Steep Spectrum (CSS) sources are powerful extragalactic radio sources with angular dimensions of the order of a few arcseconds or less. Such a compactness is apparently linked to the youth of these objects. The majority of CSSs investigated so far have been known since the early 1980s. This paper is the first in a series where we report the results of an observational campaign targeted on a completely new sample of CSSs which are significantly weaker than those investigated before. The ultimate goal of that campaign is to find out how ``weak'' CSSs compare to ``strong'', classical ones, especially with regard to the morphologies. Here we present an analysis of morphological and physical properties of five relatively large sources based on MERLIN observations at 1.6 and 5~GHz. ", + "introduction": "Compact Steep Spectrum (CSS) sources \\citep{kapahi,pw82} form a well defined class of radio sources; they are powerful, compact (projected linear sizes $\\leq$ 20~kpc and hence angular sizes of the order of a few arcseconds) and possess steep ($\\alpha\\geq 0.5$, $S\\propto\\nu^{-\\alpha}$) spectra. CSSs are identified with quasars, radio galaxies and Seyferts. \\citet{wsrps84} found a morphological separation between CSS quasars and galaxies \\citep[see also][]{spencer89,fanti90} which is similar to that observed for larger sources: radio galaxies generally have a simple double radio structure (sometimes with weak radio jets and weak radio cores) whilst quasars show either a triple structure (with a strong central component consisting of a bright jet) or complex structure. An astrophysical interpretation of the CSS phenome\\-non has been given in \\citet{fanti90}. They show that in principle small apparent sizes of CSSs could result from the projection of ``normal'' Large Symmetric Objects (LSO), however only less than 25\\% of these objects are expected to be larger sources seen close to the line of sight. Most of them are supposed to be intrinsically small objects randomly oriented on the sky and their sub-galactic apparent linear dimensions can be explained by two main hypotheses. According to the first one CSS sources are confined by the interaction of the jet with an inhomogeneous, dense and possibly turbulent medium in the host galaxy which inhibits a normal development \\citep{bmh84}. In this scenario CSSs are so called {\\it frustrated} objects. The second hypothesis \\citep{pm82, c85, mp88} --- and this one has gained more observational support recently \\citep[see e.g.][]{fanti2000} --- suggests that CSS sources may be the {\\it young} stages of future LSOs and so the compactness of these objects is just an evolutionary effect: they are small because they have not had enough time to expand to supergalactic scales. If that hypothesis is correct, CSSs can be regarded as an intermediate class between even smaller Compact Symmetric Objects (CSOs) \\citep{r96} and LSOs and, together, these three groups of radio sources make up an evolutionary sequence. One of the main argument in favour of the evolution is that CSOs and some CSS sources, namely Medium-sized Symmetric Objects (MSOs) \\citep{aug98} which are unbeamed CSSs, have similar morphologies to LSOs. On the other hand, the lobe speeds in CSO sources are high: $\\sim0.2c$ \\citep{oc98,ocp98}, so CSOs quickly evolve into larger objects and MSOs seem to be perfect candidates to become post-CSOs. Most of the CSS sources known so far have sufficiently high radio luminosity that even assuming a strong decrease in luminosity as they evolve, they remain good candidates for future large scale Fanaroff--Riley class~II (FRII) \\citep{fr74} sources with high radio luminosity. An additional support of this view comes from the morphological similarity of many CSSs to FRIIs as expected if the evolution is self-similar. All studies of CSS sources made by 1995 i.e. until publication of papers by \\citet{dffss95} and \\citet{sang95}, were based on the so called 3CRPW sample consisting of 54 sources \\citep{spencer89}. The next step in such investigations can be made through extension of the available sample of CSSs toward weaker sources. For example, \\citet{saikia01} selected a sample of 42~candidates from the S4 survey \\citep{pt78} which is complete to 0.5~Jy at 5~GHz and \\citet{fanti01} --- hereafter F2001 --- derived a new sample of 87 CSSs with flux densities $\\geq 0.8\\mathrm{Jy}$ at 408~MHz from B3-VLA survey \\citep{vig89}. Here we present a genuine method of finding weak CSSs which makes use of {\\it Faint Images of Radio Sky at Twenty} (FIRST) \\citep{wbhg97}. Its fine resolution ($5\\farcs4$) is the crucial feature for that purpose. A strong motivation to conduct the research in this direction came from the fact that a number of CSS sources weaker than those in the 3CRPW sample have already been mapped with the VLA at 8.4~GHz by \\citet{pbww92}\\footnote{In principle \\citet{pbww92} sought candidates for pointlike phase calibrators among flat-spectrum objects in Green Bank surveys but the procedure they used was different from that used by e.g. \\citet{wb92}. When we supplemented their list with flux densities from \\citeauthor{wb92} it turned out that some of their candidates actually had steep-spectra. Consequently the VLA observations revealed resolved structure and so these objects could not be used as calibrators. In this way \\citeauthor{pbww92} made a serendipitous discovery of several weak CSSs.}. It appears that those sources have angular sizes of the same range as all the CSSs known so far, yet they are significantly weaker. There are two plausible explanations: a) these sources are those CSSs we see almost exactly face-on, so their Doppler boosting is minimal, or b) these sources form a new class of ``weak'' CSSs. In 1996 we proposed observations based on newly released early results of FIRST aimed to discriminate between cases a) and b). If case a) is true then those ones with counter-jets should dominate among these new CSS sources; sources of this kind were rare in the 3CRPW sample. If case b) is true then there is a chance to find an elegant analogy: strong CSSs are ``miniature FRIIs'' whereas weak CSSs are ``miniature FRIs''. Taking into account the evolutionary scenario and assuming that radio sources evolve in a self-similar manner we might also say that strong CSSs from the 3CRPW sample evolve towards FRIIs whereas weak CSSs evolve towards FRIs. Testing such a possibility is among the goals of a series of papers resulting from interferometric (MERLIN, EVN, VLBA) observations of our FIRST-based sample of CSS candidates. \\begin{table*}[t] \\caption[]{Optical magnitudes and radio flux densities of 5 CSS sources at two frequencies} \\begin{tabular}{||c|c|c|c|r|c|c|c|c|c||} \\hline Source & 4C &RA (J2000)& DEC (J2000)& $m_R$& $m_v$& $z$& $F_{1.4\\mathrm{GHz}}$& $F_{4.85\\mathrm{GHz}}$& $\\alpha_{1.4\\mathrm{GHz}}^{4.85\\mathrm{GHz}}$\\\\ \\hline \\object{0801+303}& +30.13& 08 04 42.148& 30 12 37.91& 18.1& 19.2& 1.446& 1189& 404& 0.87\\\\ \\object{0805+406}& +40.19& 08 09 03.158& 40 32 56.72& $>20.8$& 21.1& -----& 437& 179& 0.72\\\\ \\object{0850+331}& +33.22& 08 53 21.100& 32 55 00.60& ----- & ----- & -----& 465& 208& 0.65\\\\ \\object{1201+394}& & 12 04 06.859& 39 12 18.17& 19.6& 21.4& 0.445& 468& 162& 0.85\\\\ \\object{1233+418}& & 12 35 35.706& 41 37 07.40& 17.9& 20.8& 0.25& 651& 276& 0.69\\\\ \\hline \\end{tabular} \\vspace{0.5cm} \\small{ Optical magnitudes derived from POSS plates using APM have been taken \\citet{mwhb01}.\\\\ \\object{0805+406} has not been detected at R band --- the $m_R$ value quoted is an upper limit.\\\\ The redshift of \\object{1233+418} is photometric.\\\\ Radio fluxes (in mJy) for 1400~MHz extracted from FIRST; radio fluxes (in mJy) for 4850~MHz extracted from GB6.\\\\ } \\label{table1} \\end{table*} ", + "conclusions": "MERLIN has been used to survey a sample of 60~weak Compact Steep Spectrum sources. This paper deals with five relatively large (arcsecond scale) sources. According to the evolutionary scheme compact doubles or, broadly speaking, CSOs are the progenitors of the extended doubles \\citep{pm82, c85} and the CSS sources form an evolutionary link between those most compact/youngest objects and the classical double FRIIs. All CSS sources known so far have high radio luminosities and the radio structures of those which are unbeamed have FRII structures. It seemed reasonable to suspect that the lower radio luminosity CSSs could be the progenitors of less luminous FRI objects. Our investigations of a new sample of weak CSS sources were motivated by the above-mentioned view. The triple or double structures of the five CSS sources and the presence of one-sided core--jet structures indicate they are more similar to FRII objects than FRIs. The radio structure of \\object{1201+394} is also similar to the structure of FRII object because of edge-brightening of the lobes. The remaining four sources seem to be moderately beamed. Their structures consist of a core and one-sided jet. None of these sources have counter-jets. This means that the low luminosities of these sources are not a consequence of a lack of significant Doppler boosting. We claim, therefore, they constitute a new class of ``weak'' CSS sources." + }, + "0112/hep-ph0112212_arXiv.txt": { + "abstract": "We have recently proposed a mechanism of photon-axion oscillations as a way of rendering supernovae dimmer without cosmic acceleration. Subsequently, it has been argued that the intergalactic plasma may interfere adversely with this mechanism by rendering the oscillations energy dependent. Here we show that this energy dependence is extremely sensitive to the precise value of the free electron density in the Universe. Decreasing the electron density by only a factor of 4 is already sufficient to bring the energy dependence within the experimental bounds. Models of the intergalactic medium show that for redshifts $z<1$ about 97\\% of the total volume of space is filled with regions of density significantly lower than the average density. From these models we estimate that the average electron density in most of space is lower by at least a factor of 15 compared to the estimate based on one half of all baryons being uniformly distributed and ionized. Therefore the energy dependence of the photon-axion oscillations is consistent with experiment, and the oscillation model remains a viable alternative to the accelerating Universe for explaining the supernova observations. Furthermore, the electron density does give rise to a sufficiently large plasma frequency which cuts off the photon-axion mixing above microwave frequencies, shielding the cosmic microwave photons from axion conversions and significantly relaxing the lower bounds on the axion mass implied by the oscillation model. ", + "introduction": " ", + "conclusions": "" + }, + "0112/astro-ph0112277_arXiv.txt": { + "abstract": "Despite the fact that the physics of the cosmic microwave background anisotropies is most naturally expressed in Fourier space, pixelised maps are almost always used in the analysis and simulation of microwave data. A complementary approach is investigated here, in which maps are used only in the visualisation of the data, and the temperature anisotropies and polarization are only ever expressed in terms of their spherical multipoles. This approach has a number of advantages: there is no information loss (assuming a band-limited observation); deconvolution of asymmetric beam profiles and the temporal response of the instrument are naturally included; correlated noise can easily be taken into account, removing the need for additional `destriping'; polarization is also analysed in the same framework; and reliable estimates of the spherical multipoles of the sky and their errors are obtained directly for subsequent component separation and power spectrum estimation. The formalism required to analyse experiments which survey the full sky by scanning on circles is derived here, with particular emphasis on the \\emph{Planck} mission. A number of analytical results are obtained in the limit of simple scanning strategies. Although there are non-trivial computational obstacles to be overcome before the techniques described here can be implemented at high resolution, if these can be overcome the method should allow for a more robust return from the next generation of full-sky microwave background experiments. ", + "introduction": "\\label{sec:intro} The cosmic microwave background (CMB) represents the single most powerful probe of the early universe available to modern astronomy. The pioneering results of the \\emph{Cosmic Background Explorer} (\\emph{COBE}; Smoot et al.\\ 1992) are being extended by the current generation of ground-based and balloon-borne experiments (e.g.\\ Scott et al.\\ 1996; Tanaka et al.\\ 1996; Netterfield et al.\\ 1997; de Oliveira-Costa et al.\\ 1998; Coble et al.\\ 1999; de Bernardis et al.\\ 2000; Hanany et al.\\ 2000; Wilson et al.\\ 2000; Padin et al.\\ 2001), but it is the upcoming satellite experiments that are set to revolutionise the field. By the end of 2002 the \\emph{Microwave Anisotropy Probe}\\footnote{http://map.gsfc.nasa.gov/} (\\emph{MAP}) will have observed the entire sky with a resolution of $\\ge 0\\fdg 21$ and \\emph{Planck}\\footnote{http://astro.estec.esa.nl/Planck/}, scheduled for launch in 2007, should have the capacity to produce high sensitivity, all-sky maps to a resolution of $\\sim$ 0\\fdg 1 in four of its ten frequency channels. However, it is not sufficient merely to obtain such extraordinary measurements; careful analysis of the data-sets is critical if these missions are to fulfill their promised scientific goals. Most present techniques for CMB data analysis employ pixelised maps of the sky. Both Bond et al.~\\shortcite{bond99} and Borrill~\\shortcite{borrill99} describe methods to create maximum-likelihood maps from time-ordered data and telescope pointing information, and these techniques have been successfully applied in the analysis of the BOOMERanG~\\cite{deb00,net01} and MAXIMA~\\cite{hanany00,lee01} observations. From the maps and their uncertainties, the power spectrum of the microwave sky can be estimated in a number of different ways (Bond, Jaffe \\& Knox 1998; Borrill 1999; Oh, Spergel \\& Hinshaw 1999; Wandelt, Hivon \\& G\\'{o}rski 2001), again employing maximum-likelihood techniques. These methods have a number of useful features: a map represents vast data compression relative to the time-ordered data (e.g.\\ Borrill 1999), but is still a sufficient statistic for cosmological parameter estimation; there are usually negligible pixel-pixel noise correlations in the beam-smoothed map; maps provide important `reality checks', and can be inspected by eye; and unobserved or contaminated regions of the sky can be simply removed from subsequent analysis by excluding the associated pixels (e.g.\\ Oh et al.\\ 1999). From a more practical point of view, the conventional map-making method is `tried and tested' now. However, the use of pixelised maps during the data analysis also involves some compromises: the real microwave sky does not consist of a number of regions of uniform temperature (the prior hypothesis used to make a map), and so information is lost in map-making; subsequent steps in the analysis pipeline, such as component separation, cannot easily be performed efficiently in real space~\\cite{hobson98}; there is no single optimal or obvious choice of pixelisation scheme; and deconvolution of the (asymmetric) beam profile and temporal response of instrument, and removal of scan-synchronous instrument effects require awkward additional processing during map-making~\\cite{del98a,revenu00,wu01}. Another potential problem that has received only limited attention thus-far (e.g. Tegmark \\& de Oliveira-Costa 2001) is how to create maps from polarization data (i.e.\\ a map for each of the Stokes parameters, or their gradient and curl components) -- there is currently no robust algorithm for the treatment of several polarization-specific systematic effects that plague CMB polarimetry experiments. It is with these points in mind that a complementary, harmonic method of data analysis is presented here. For full-sky surveys the data analysis process would proceed from the time-ordered data directly to the set of spherical multipole moments, which would take the place of a real-space map. For band-limited observations (effectively ensured by the finite experimental beam), this harmonic reconstruction of the sky involves essentially no loss of information. The separation of the various astrophysical components of the microwave sky could then be performed quite naturally in multipole space (e.g.\\ Hobson et al.\\ 1998), and, if necessary, the Galactic plane could be removed while retaining orthonormality of the basis set (G\\'{o}rski 1994; Mortlock, Challinor \\& Hobson 2001; see also Section~\\ref{subsec:power}). Power spectrum estimation could then proceed naturally from this point; for instance the efficient method presented by Oh et al.~\\shortcite{oh99} could be used, although some generalisation would be required to remove the dependence on forward and inverse fast spherical transforms, which Oh et al.\\ use to apply the inverse noise matrix efficiently in the space of multipoles and to reduce memory requirements. The multipole moments of the polarized components could be obtained using an almost identical formalism; obviously it would incur an additional overhead for each component to be reconstructed, but no further conceptual development would be required. We have endeavoured to present the harmonic method as an end-to-end solution for the processing of time-ordered data into power spectra. For such an analysis pipeline the benefits of clear error propagation from timelines to power spectra afforded by our harmonic methods are maximised. However, we do not pretend that all aspects of CMB analysis are best performed in the harmonic domain; instead we view harmonic methods as being complementary to standard map-based techniques. An obvious example where map-based processing is clearly desirable is the analysis of localised foregrounds (or CMB features). In addition, the inversion from observations over only a fraction of the sky to the spherical multipoles rapidly becomes singular as the resolution of the observations is increased. Although it may be possible to regularise the inversion in such cases (e.g.\\ with prior information on the power spectrum), or to adopt a basis set adapted to the observed patch, the processing we describe in this paper is best considered in the context of full-sky observations. Even then, some aspects of the processing, such as cutting out contaminated regions close to the galactic plane for power spectrum estimation, require rather cumbersome operations if we insist on working directly with the full sky spherical multipoles rather than a map synthesised from them. The principles of harmonic data analysis as described above are quite general, and could be applied to any full-sky CMB observations. However, it is particularly suited to experiments with a circular scanning strategy, such as the \\emph{Planck} mission. The data are obtained by a combination of rotations (i.e.\\ the motion of the detector across the sky) and convolutions (of the beam with the true microwave sky, and this subsequent signal with the temporal response of the instrument), and these operations are most simply expressed in the harmonic basis (Delabrouille, G\\'{o}rski \\& Hivon 1998a; Challinor et al.\\ 2000; Wandelt \\& G\\'{o}rski 2001). Hence we propose that the time-ordered data be transformed to the Fourier basis at the earliest opportunity: van Leeuwen et al.~\\shortcite{van01} describe how to construct point source catalogues and calibrated ring-sets simultaneously, in which the data around each ring is represented in Fourier space. Some aspects of this process are necessarily mission-specific, and, for the sake of generality, it is assumed here that it is possible to obtain calibrated ring-sets (and their errors) in this form. The focus of this paper is on the subsequent `map'-making: combining the ring-sets to obtain an estimate of the spherical multipoles for the temperature anisotropies and polarization. In this respect, we build on earlier work by Delabrouille et al.~\\shortcite{del98b}, who analysed the statistics of the power spectrum of the Fourier modes on a single ring, and the relation with the power spectrum of the underlying temperature field on the sphere. More recently, Wandelt \\& Hansen~\\shortcite{wan01a} have proposed the `ring-torus' method for estimating the temperature power spectrum from the two-dimensional Fourier transform of data obtained on a set of rings with centres equally spaced on a small circle. While the general philosophy of the `ring-torus' method coincides with that advocated here, the methodology presented here is somewhat more general. By only making use of one-dimensional Fourier data, we are able to deal with arbitrary ring-sets and to include instrument effects such as the inevitable variations in the scanning properties of the telescope~\\cite{van01}. Furthermore, since we do not proceed directly to the power spectrum, we can make full use of Fourier-based component separation algorithms (e.g.\\ Hobson et al.\\ 1998) to deal carefully with foreground contamination. The cost of maintaining this generality is the increased computational requirements of our methods over those that exploit special symmetries of the survey. The general formalism is described in Section~\\ref{sec:hdm}, which proceeds from a model of the instrument to the maximum-likelihood solution for the spherical multipoles of the temperature anisotropy and polarization, and their associated errors. The structure of the error covariance matrix, including polarization, is studied for some simple scan strategies in Section~\\ref{sec:scan}. Making use of some results given in Appendix B, we are able to make contact with several well-known analytic results for idealised experiments, previously obtained from arguments at the level of the map. In Section~\\ref{sec:disc} we discuss a number of important issues that arise during the map-making phase, including correcting for scan-synchronous instrument effects, the overall calibration of the instrument to external standards, and the treatment of low frequency noise. The latter discussion includes a novel method for dealing with uncertainties in the low frequency noise power spectrum or insensitivity of the experiment to the monopole. Section~\\ref{sec:sub} reviews the subsequent processing of the frequency maps, including component separation and power spectrum estimation, within the context of the harmonic data model. Finally, we conclude in Section~\\ref{sec:conc} by reviewing the relative merits of the harmonic method, and suggest directions for future development. ", + "conclusions": "\\label{sec:conc} Most current techniques for analysing CMB data on the sphere are based on the use of pixelised maps. The basis for a complementary approach, based on spherical harmonic coefficients of the intensity and polarization, has been derived here for the case of experiments that scan the full sky in circles. Our method offers a number of advantages, most notably the refined treatment of non-ideal beam effects, the ability to handle correlated (low frequency) noise in an optimal, but efficient, manner, and the seamless way in which polarized data can be analysed alongside total intensity data. In principle, harmonic methods can be used to develop an end-to-end pipeline for the analysis of full-sky survey data, allowing the clean propagation of noise and other errors from the time-ordered data to the CMB power spectra. In practice, the methods described here are best regarded as being complementary to more standard map-based techniques rather than a replacement. Undoubtly, there are analysis projects that are better suited to map-based techniques -- typically those concerned with the science of local features in the maps, such as foregrounds. In addition, power spectrum estimation becomes cumbersome with purely harmonic methods if we have reason to question the foreground contamination of certain linear combinations of the spherical multipoles (e.g.\\ those corresponding to features localised in the galactic plane). The biggest hurdle facing a practical implementation of the harmonic `map'-making method at the resolution demanded by the upcoming satellite missions, is the need to solve the $O(l_{{\\rmn{max}}}^2)\\times O(l_{{\\rmn{max}}}^2)$ linear system in equation~(\\ref{eq:28}). Unlike conventional map-making in the presence of noise correlations, the problem lies not in the inversion of the noise covariance matrix, since we are working in a representation where this matrix is already very sparse. The inversion of ${\\mathbfss{A}}^\\dagger {\\mathbfss{N}}^{-1} {\\mathbfss{A}}$ can be avoided by adopting iterative techniques with a block-diagonal preconditioner. The main computational overhead arises from computing and storing the large, non-sparse matrix $\\mathbfss{A}$ and applying it to the sky multipoles and its transpose to the Fourier ring data. These operations can be performed very efficiently using fast Fourier transform techniques for the rather idealised case of constant latitude scanning, with rings uniformally spaced in azimuth and no variations in instrument properties during the mission~\\cite{wan01b}. Unfortunately, it does not appear that such techniques can be easily extended to more realistic scan strategies. An assessment of some of these numerical problems, together with a number of potential solutions, will be given in Mortlock et al.\\ (in preparation)." + }, + "0112/astro-ph0112088_arXiv.txt": { + "abstract": "We describe results from a fully self--consistent three dimensional hydrodynamical simulation of the formation of one of the first stars in the Universe. Dark matter dominated pre-galactic objects form because of gravitational instability from small initidal density perturbations. As they assemble via hierarchical merging, primordial gas cools through ro-vibrational lines of hydrogen molecules and sinks to the center of the dark matter potential well. The high redshift analog of a molecular cloud is formed. When the dense, central parts of the cold gas cloud become self-gravitating, a dense core of $\\sim 100\\Ms$ undergoes rapid contraction. At densities $n>10^9 \\ccc$ a $1\\Ms$ proto-stellar core becomes fully molecular due to three--body \\HH formation. Contrary to analytical expectations this process does not lead to renewed fragmentation and only one star is formed. The calculation is stopped when optical depth effects become important, leaving the final mass of the fully formed star somewhat uncertain. At this stage the protostar is acreting material very rapidly ($\\sim 10^{-2}\\Ms \\yr^{-1}$). Radiative feedback from the star will not only halt its growth but also inhibit the formation of other stars in the same pre--galactic object (at least until the first star ends its life, presumably as a supernova). We conclude that at most one massive ($M\\gg1\\Ms$) metal free star forms per pre--galactic halo, consistent with recent abundance measurements of metal poor galactic halo stars. ", + "introduction": "Chemical elements heavier than Lithium are synthesized in stars. Such ``metals'' are observed at times when the Universe was only $\\lsim 10$\\% of its current age in the inter--galactic medium (IGM) as absorption lines in quasar spectra (see Ellison et al. 2000, and references therein). Hence, these heavy elements not only had to be synthesized but also released and distributed in the IGM within the first billion years. Only supernovae of sufficiently short lived massive stars are known to provide such an enrichment mechanism. This leads to the prediction that {\\it the first generation of cosmic structures formed massive stars (although not necessarily only massive stars).} In the past 30 years it has been argued that the first cosmological objects form globular clusters (\\r{PD68}), super--massive black holes (\\r{H69}), or even low mass stars (\\r{PSS83}). This disagreement of theoretical studies might at first seem surprising. However, the first objects form via the gravitational collapse of a thermally unstable reactive medium, inhibiting conclusive analytical calculations. The problem is particularly acute because the evolution of all other cosmological objects (and in particular the larger galaxies that follow) will depend on the evolution of the first stars. Nevertheless, in comparison to present day star formation, the physics of the formation of the first star in the universe is rather simple. In particular: \\begin{itemize} \\item the chemical and radiative of processes in the primordial gas are readily understood. \\item strong magnetic fields are not expected to exist at early times. \\item by definition no other stars exist to influence the environment through radiation, winds, supernovae, etc. \\item the emerging standard model for structure formation provides appropriate initial conditions. \\end{itemize} In previous work we have presented three--dimensional cosmological simulations of the formation of the first objects in the universe (\\r{A95}, \\r{AANZ}) including first applications of adaptive mesh refinement (AMR) cosmological hydrodynamical simulations to first structure formation (\\r{ABN99}, \\r{ABN00}, ABN hereafter) . In these studies we achieved a dynamic range of up to $2\\tento{5}$ and could follow in detail the formation of the first dense cooling region far within a pre--galactic object that formed self--consistently from linear density fluctuation in a cold dark matter cosmology. Here we report results from simulations that extend our previous work by another 5 orders of magnitude in dynamic range. For the first time it is possible to bridge the wide range between cosmological and stellar scale. ", + "conclusions": "Previously we discussed the formation of the pre--galactic object and the primordial ``molecular cloud'' that hosts the formation of the first star in the simulated patch of the universe (\\r{ABN00}). These simulations had a dynamic range of $\\sim 10^5$ and identified a $\\sim 100\\Ms$ core within the primordial ``molecular cloud'' undergoing renewed gravitational collapse. The fate of this core was unclear because there was the potential caveat that three body \\HH\\ formation could have caused fragmentation. Indeed this further fragmentation had been suggest by analytic work (\\r{Silk83}) and single zone models (\\r{PSS83}). The three dimensional simulations described here were designed to be able to test whether the three body process will lead to a break up of the core. {\\sl No fragmentation due to three body \\HH\\ formation is found.} This is to a large part because of the slow quasi--hydrostatic contraction found in ABN which allows sub--sonic damping of density perturbations and yields a smooth distribution at the time when three body \\HH\\ formation becomes important. Instead of fragmentation a single fully molecular proto--star of $\\sim 1\\Ms$ is formed at the center of the $\\sim 100\\Ms$ core. However, even with extraordinary resolution, the {\\it final} mass of the first star remains unclear. Whether all the available cooled material of the surroundings will accrete onto the proto--star or feedback from the forming star will limit the further accretion and hence its own growth is difficult to compute in detail. Within $10^4\\yr$ about $70\\Ms$ may be accreted assuming that angular momentum will not slow the collapse (Fig.~\\ref{acrete}). The maximum of the accretion time of $\\sim 5\\tento{6}\\yr$ is at $\\sim 600 \\Ms$. However, stars larger than $100\\Ms$ will explode within $\\sim 2\\Myr$. Therefore, it seems unlikely (even in the absence of angular momentum) that there would be sufficient time to accrete such large masses. \\vspace{.1cm} \\begin{figurehere} \\centerline{\\psfig{file=abelfig5.eps,width=9cm}} \\caption{The accretion time as function of enclosed gas mass. The line with symbols gives $M(r)/[4 \\pi \\rho(r) r^2 |v_r(r)| ]$. The solid line simply shows how long it would take the mass to move to $r=0$ if it were to to keep its current radial velocity ($r/v_r(M)$). }\\label{acrete} \\end{figurehere} \\vspace{.2cm} A one solar mass proto--star will evolve too slowly to halt substantial accretion. From the accretion time profile (Fig.~\\ref{acrete}) one may argue that a more realistic minimum mass limit of the first star should be $\\gsim 30 \\Ms$ because this amount would be accreted within a few thousand years. This is a very short time compared to expected proto--stellar evolution times. However, some properties of the primordial gas may make it easier to halt the accretion. One possibility is the destruction of the cooling agent, molecular hydrogen, without which the acreting material may reach hydrostatic equilibrium. This may or may not be sufficient to halt the accretion. One may also imagine that the central material heats up to $10^4$ K, allowing Lyman-$\\alpha$ cooling from neutral hydrogen. That cooling region may expand rapidly as the internal pressure drops because of infall, possibly allowing the envelope to accrete even without molecular hydrogen as cooling agent. Additionally, radiation pressure from ionizing photons as well as atomic hydrogen Lyman series photons may become significant and eventually reverse the flow. The mechanisms discussed by Haehnelt (1993) on galactic scales will play an important role for the continued accretion onto the proto--star. This is an interplay of many complex physical processes because one has a hot ionized Str\\\"omgren sphere through which cool and dense material is trying to accrete. In such a situation one expects a Raleigh--Taylor type instability that is modified via the geometry of the radiation field. At the final output time presented here there are $\\sim 4\\tento{57}$ hydrogen molecules in the entire protogalaxy. Also the \\HH\\ formation time scale is long because there are no dust grains and the free electrons (needed as a catalyst) have almost fully recombined. Hence, as soon as the the first UV photons of Lyman Werner band frequencies are produced there will be a rapidly expanding photo--dissociating region (PDR) inhibiting further cooling within it. This photo-dissociation will prevent further fragmentation at the molecular cloud scale. I.e. no other star can be formed within the same halo before the first star dies in a supernova. The latter, however, may have sufficient energy to unbind the entire gas content of the small pre--galactic object it formed in (\\r{MF99}). This may have interesting feedback consequences for the dispersal of metals, entropy and magnetic field into the intergalactic medium (\\r{F98}, \\r{CB00}). Smoothed particle hydrodynamics (SPH, e.g. \\r{M92}), used extensively in cosmological hydrodynamics, has been employed (\\r{BCL99}) to follow the collapse of solid body rotating uniform spheres. The assumption of coherent rotation causes these clouds to collapse into a disk which developes filamentary structures which eventually fragment to form dense clumps of masses between $100$ and $1000$ solar masses. It has been argued that these clumps will continue to accrete and merge and eventually form very massive stars. These SPH simulation have unrealistic initial conditions and much less resolution then our calculations. However, they also show that many details of the collapse forming a primordial star are determined by the properties of the hydrogen molecule. We have also simulated different initial density fields for a Lambda CDM cosmology. There we have focused on halos with different clustering environments. Although we have not followed the collapse in these halos to proto-stellar densities, we have found no qualitative differences in the ``primordial molecular cloud'' formation process as discussed in ABN. Also other AMR simulations (\\r{MBA01}) give consistent results on scales larger than $1\\pc$. In all cases a cooling flow forms the primordial molecular cloud at the center of the dark matter halo. We conclude that the molecular cloud formation process seems to be independent of the halo clustering properties and the adopted CDM type cosmology. Also the mass scales for the core and the proto--star are determined by the local Bonnor--Ebert mass. Consequently, we expect the key results discussed here to be insensitive to variations in cosmology or halo clustering." + }, + "0112/astro-ph0112041_arXiv.txt": { + "abstract": "The MIT and CXC ACIS teams have explored a number of measures to ameliorate the effects of radiation damage suffered by the ACIS FI CCDs. One of these measures is a novel CCD read-out method called ``squeegee mode''. A variety of different implementations of the squeegee mode have now been tested on the I0 CCD. Our results for the fitted FWHM at Al-K$\\alpha$ and Mn-K$\\alpha$ clearly demonstrate that all the squeegee modes provide improved performance in terms of reducing CTI and improving spectral resolution. Our analysis of the detection efficiency shows that the so-called squeegee modes ``Vanilla'' and ``Maximum Observing Efficiency'' provide the same detection efficiency as the standard clocking, once the decay in the intensity of the radioactive source has been taken into account. The squeegee modes which utilize the slow parallel transfer (``Maximum Spectral Resolution'', ``Maximum Angular Resolution'', and ``Sub-Array'') show a significantly lower detection efficiency than the standard clocking. The slow parallel transfer squeegee modes exhibit severe grade migration from flight grade 0 to flight grade 64 and a smaller migration into ASCA g7. The latter effect can explain some of the drop in detection efficiency. There are a few observational penalities to consider in using a squeegee mode. Utilizing any squeegee mode causes a loss of FOV near the aimpoint (4 to 16$''$ strips along the full length of the CCDs), as well as the attendant dead-time increase. Secondly, the cost of the software implementation and its testing will be significant. Lastly, each squeegee mode ``flavor'' would require lengthy, mode-specific calibration observations. Therefore, since an efficacious, ground-based CTI corrector algorithm is now available (see paper by Plucinsky, Townsley, \\textit{et al.} in this proceedings), a scientific judgment will have to be made to determine which, if any, squeegee modes should be developed and calibrated for use by Chandra observers. ", + "introduction": "The ACIS Operations team, combining elements from the CXC and the ACIS MIT$/$IPI team, has been developing and testing the new squeegee modes since the Spring of 2000. In squeegee mode, charge is collected in the top few rows of the CCD and then swept across the imaging array once per readout, thus filling some of the radiation-induced electron traps that cause degraded performance. The development of a novel method of reading out the ACIS CCDs was first developed and tested on ground CCDs that are similar to flight CCDs (see Prigozhin, \\textit{et al.} 2000 for a characterization of the radiation damage of ACIS CCDs). This new mode ameliorates some of the effects caused by the radiation damage suffered early in the mission by the FI CCDs. For a more comprehensive discussion on the design and clocking method of squeegee mode, see Bautz and Kissel (2000). A description of the different squeegee modes is included in an internal MIT$/$ACIS Memo by Bautz and Grant (2000). That memo also presents an analysis of the squeegee modes and describes the trade-offs in choosing between the various squeegee modes. In this paper, we present the results of the various squeegee measurements on the I0 CCD. \\begin{table}[tbp] \\centering \\caption[ ]{{\\bf List of I0 Squeegee Tests } } \\label{datasets} \\begin{tabular}{|l|l|l|l|l|} \\hline Test & OBSID & Date & Exp(s) & Description \\\\ \\hline \\hline Control & 62895 & 20 Feb 2000 & 9,023.4 & I0, 3.2s, standard clocking\\\\ & & & & ($40\\mu$s par xfr, no 2x2 sum)\\\\ & & & & ``Control'' non-Squeegee \\\\ \\hline L1 & 62042 & 30 May 2000 & 8,225.5 & I0, 3.3s, 16 row sq, 32 row exc win, no 2x2 sum, \\\\ & & & & $40\\mu$s par xfr, 24 flushes, 1010 rev clks \\\\ & & & & ``Vanilla Squeegee'' \\\\ \\hline L2 & 62019 & 21 Jun 2000 & 8,623.6 & I0, 1.9s, 16 row sq, 24 row exc window, 2x2 sum, \\\\ & & & & $320\\mu$s par xfr, 24 flushes, 1026+32 rev clks \\\\ & & & & Maximum Spectral Resolution \\\\ \\hline L4 & 62007 & 01 Jul 2000 & 8,517.3 & I0, 3.3s, 2 row sq, 8 row exc win, no 2x2 sum, \\\\ & & & & $320\\mu$s par xfr, 24 flushes, 1026+32 rev clks \\\\ & & & & Maximum Angular Resolution \\\\ \\hline L7 & 61981 & 30 Jul 2000 & 8,551.5 & I0, 1.8s, 2 row sq, 8 row exc win, 2x2 sum, \\\\ & & & & $40\\mu$s par xfr, 24 flushes, 1026+32 rev clks \\\\ & & & & Maximum Observing Efficiency \\\\ \\hline L9 & 61949 & 30 Aug 2000 & 8,356.9 & I0, 1.2s, 2 row sq, 8 row exc win, no 2x2 sum, \\\\ & & & & $320\\mu$s par xfr, 24 flushes, 1026+32 rev clks, \\\\ & & & & 256 row sub-array \\\\ \\hline \\end{tabular} \\end{table} This analysis utilized the standard level 0 and level 1 data products produced by the CXC Data System. The datasets included in this analysis are listed in Table 1. The control run (OBSID 62895) was a long charge transfer inefficiency (CTI) measurement. There are five ``flavors'' of Squeegee discussed in this paper: L1 ``Vanilla'', L2 ``Maximum Spectral Resolution'', L4 ``Maximum Angular Resolution'', L7 ``Maximum Observing Efficiency'', and L9 ``Sub-Array near the aimpoint''. Table 1 summarizes the important parameters which distinguish one squeegee mode from another: static integration time, number of squeegee rows (2 or 16), number of rows in the exclusion window (8 or 24 or 32), on-chip summing (yes or no), parallel transfer time ($40\\mu$s or $320\\mu$s), number of frame flushes (always 24 for these squeegees), and number of reverse clocks (1010 or 1026+32). In addition, the table includes the average exposure time for these measurements. With all of these squeegee modes, the exposure time varies from row-to-row. The number listed is the correct exposure time for the middle of the CCD or for the middle of the sub-array for squeegee L9. In the analysis presented in this paper, the row-to-row variation in exposure time has been included by computing the exposure for each of the 32 row elements. ", + "conclusions": "We confirm that all of the tested squeegee modes improve the spectral resolution of the I0 CCD compared to the standard clocking. Our analysis of the detection uniformity indicates that squeegees L1 and L7 have the same detection efficiency as the standard clocking after correcting for the decay in the intensity of the radioactive source, while the squeegees L2, L4, and L9 still exhibit a lower detection efficiency. The discrepancy is as large as 20\\% for squeegee L4 at the top of the CCD. L2 and L4 also produce a highly spatially-dependent grade distribution. We suggest that the slow parallel transfer of both these modes is the likely explanation for this effect. We suggest that this effect should be investigated further with the hope that a squeegee mode can be developed which optimizes the spectral resolution, the detection efficiency, and the {\\em uniformity} of the detection efficiency." + }, + "0112/astro-ph0112331_arXiv.txt": { + "abstract": "A comprehensive study of the measurement of star formation histories from colour-magnitude diagrams (CMDs) is presented, with an emphasis on a variety of subtle issues involved in the generation of model CMDs and maximum likelihood solution. Among these are the need for a complete sampling of the synthetic CMD, the use of of proper statistics for dealing with Poisson-distributed data (and a demonstration of why $\\chi^2$ must not be used), measuring full uncertainties in all reported parameters, quantifying the goodness-of-fit, and questions of binning the CMD and incorporating outside information. Several example star formation history measurements are given. Two examples involve synthetic data, in which the input and recovered parameters can be compared to locate possible flaws in the methodology (none were apparent) and measure the accuracy with which ages, metallicities, and star formation rates can be recovered. Solutions of the histories of seven Galactic dwarf spheroidal companions (Carina, Draco, Leo I, Leo II, Sagittarius, Sculptor, and Ursa Minor) illustrate the ability to measure star formation histories given a variety conditions -- numbers of stars, complexity of star formation history, and amount of foreground contamination. Significant measurements of ancient $> 8$ Gyr star formation are made in all seven galaxies. Sculptor, Draco, and Ursa Minor appear entirely ancient, while the other systems show varying amounts of younger stars. ", + "introduction": "The evolution of galaxies can be studied two ways -- one can either look at high redshift to observe the past directly, or one can look at the fossil remains of past events in nearby galaxies. These approaches are complementary, as the first is more direct (the ancient light is being observed now) but allows only a statistical comparison of the events seen happening at different ages (we cannot be entirely sure which systems at high redshift are analogous to which systems in the nearby universe). In contrast the measurement of the star formation history of a nearby galaxy (one whose stellar content is resolved) allows one to trace the history of a single system, but it is difficult to determine that history in an unambiguous way. The measurement of star formation histories via comparisons of observed and synthetic colour-magnitude diagrams (CMDs) is an active field that is evolving rapidly. The first papers on the topic arrived in the literature only slightly more than a decade ago, with Tosi et al. (1989) and Bertelli et al. (1992) two early attempts to derive the star formation histories of composite populations (stars of a range of ages and metallicities). As opposed to isochrone fitting in single-population objects, such as globular clusters, the measurement of a star formation history of a composite system is a daunting task -- SFR(t,Z), distance, extinction/reddening, initial mass function (IMF), and binary distribution are all unknowns at some level; while the comparison of CMDs was done subjectively. In order to cope with the vast combinations of parameters possible given limitations in computer speed and the number of subjective comparisons that could be made, these early studies limited the parameter space, generally measuring only a small set of SFR(t) functions and assuming fixed values for all other parameters. The work of Gallart et al. (1996a), studying the old stellar content of NGC 6822, was the first attempt to quantify the subjective CMD comparisons. In that work, the authors constructed a large number of parameters, each of which measured the position, size, and/or number of stars of a certain feature of the CMD. This allowed the first quantitative judgment of a star formation history, although both the procedures for generating synthetic CMDs and comparing CMDs were still extremely slow, forcing a solution of only SFR(t) and Z(t). The shift to a fully quantitative analysis was proposed independently by Dolphin (1997) and Aparicio, Gallart, \\& Bertelli (1997), who proposed binning the CMD into sections and performing a $\\chi^2$ minimization on the number of stars in each section to determine the star formation history. Dolphin (1997) demonstrated the sensitivity of such a method to metallicity, distance, and extinction, and as well as its ability to correctly reconstruct the star formation history of a synthetic population; Aparicio et al. (1997) applied a remarkably similar algorithm to a study of LGS 3. The advantage in such a technique lay in its ability to use all parts of the CMD in measuring the star formation history -- thus allowing it to be used on photometry of any quality and depth -- as well as the obvious advantage of having a single-parameter fit that can be used in a numerical minimization. The number of groups working on this topic continues to increase; Tolstoy \\& Saha (1996); Holtzman et al. (1999); Olsen (1999); Hernandez, Gilmore, \\& Valls-Gabaud (2000); and Harris \\& Zaritsky (2001) is only a partial list of other groups that are using CMDs to measure past star formation histories. These techniques have been applied to many of the Local Group galaxies, as well as a few galaxies just outside the Local Group. Despite the large number of papers on this topic, the literature lacks thorough methodology papers describing modern techniques, largely because of the incremental improvements in methods that have been implemented by each group. An example of this is the series of papers by Dolphin -- Dolphin (1997) presenting the initial method; Dolphin (2000a), Dolphin et al. (2001a), Miller et al. (2001), and Dolphin (2001) each containing minor improvements to the technique -- which forces the reader to follow a paper trail to determine what any one group is currently doing. Another significant void in the literature is a realistic estimation of how well one can measure star formation histories under a variety of conditions -- number of stars in the CMD, amount of foreground contamination, and complexity of star formation. The present work attempts to fill these needs in the literature, in addition to addressing commonly-made mistakes. This paper is divided into two main sections -- a detailed description of how to measure star formation histories from CMDs and application to artificial and real data. ", + "conclusions": "An examination of the technique for measuring star formation histories has been presented. While the underlying concept -- finding the star formation history most likely to have produced the observed data -- is straightforward, there are a number of potential traps that must be overcome. In generating synthetic CMDs, one must take care to ensure that all possible outcomes have been sampled; this is not done by ``random drawing'' techniques in which a certain number of stars are randomly drawn and placed on the CMD. Instead, it is necessary to make a true model CMD -- a CMD that represents the probability distribution from which the data could have been drawn. The first step is to make fine interpolations of the isochrones in age, metallicity, and mass so that all possible single stars are accounted for. One must also account for the possibility of binaries by considering a number of possible secondary passes sufficiently large to create a smooth model CMD. The final step in generating a partial CMD (CMD for a small range of age and metallicity) is to apply the results of artificial star tests. The model CMD can be generated from any combination of the partial CMDs (different combinations correspond to different star formation histories), plus a model of foreground contamination and models of bad detections. I have demonstrated the inadequacy of a $\\chi^2$ minimization when fitting Poisson-distributed data (as is the case here). Specifically, a $\\chi^2$ minimization will always minimize with the wrong star formation history; the only question is how wrong the answer will be. Instead, a Poisson likelihood ratio is recommended, the equation given in equation \\ref{eq_plr}. It has also been demonstrated that the ``Saha $W$'' (Saha 1998) statistic is not designed for model-data comparisons. The Bayesian inference scheme of Tolstoy \\& Saha (1996) provides an accurate solution of relative star formation rates but not the overall mean star formation rate. The question of binning vs. non-binning is demonstrated to be unimportant, as the same star formation rate will be obtained so long as the bin sizes are as small as the smallest features of the model CMD. Finally, techniques for measuring uncertainties and determining the overall fit quality are given, as well as a method in which outside data (such as a red giant metallicity distribution) can be incorporated into the fit without use of a prior. The technique was then applied to a pair of synthetic galaxies -- one single-population and one composite-population. The star formation history of the single-population system was measured with an age accuracy of $\\pm 0.03$ dex and distance and extinction accuracy of 0.02 magnitudes, provided that at least the RGB and HB were included in the data (depth of $M_V = +2$). Most of the constraints were lost, however, when restricting the solution to only the upper RGB (depth of $M_V = 0$); this introduced an age uncertainty of $\\pm 0.2$ dex into the solution. Although the quality of the fit was severely degraded when using an intentionally-wrong set of age bins, we note that the measured distance, extinction, and star formation history were all correct. The star formation history of the synthetic composite-population system was measured with less accuracy, with resolution of roughly $\\pm 0.07$ dex producing reasonable signal-to-noise with photometric depth of $M_V = +2$. However, the solution with a photometric limit of $M_V = 0$ was again very uncertain, with age resolution degraded to $\\pm 0.25$ dex. The quoted resolutions, of course, are dependent upon the number of stars in the observed field; the uncertainties scale as $1/\\sqrt{N}$. Finally, I showed measurements of the star formation histories of seven dwarf spheroidal companions. While each data set had a different quality (number of stars, photometric depth, and amount of foreground contamination), the ability to accurately measure uncertainties allows one to give the best answer and the uncertainty in the measurement for each object. Thus the star formation history can always be measured -- even with the very poor Sculptor CMD -- but better data will naturally result in smaller uncertainties. The technique-related findings of this study can be summarized as follows: \\begin{enumerate} \\item In every case, the calculated star formation history matched with the qualitative star formation history obtained by a cursory examination of the CMD. In nearly every case, the distance and extinction were consistent with literature values. \\item The number of stars with $M_V < +4$ required to produce results with signal-to-noise $>$ 1 at moderate resolution appears to be about 150 for an old system and 500-1000 for a system with many young stars. \\item Even with the uncertainties in the isochrones, all CMDs were well-fit. The largest $\\chi^2_{eff}$ was 1.16, and only the Leo I fit was worse than 2.5 $\\sigma$ from an ideal solution. \\item In the case of the Sagittarius dwarf, a large amount of foreground contamination (more foreground stars than Sagittarius stars) does not add significantly to the fit uncertainties. This is likely because the main sequence and MSTO of Sagittarius are sufficiently separated from the bulge main sequence. \\end{enumerate} Scientifically, the results are limited by the fact that only a small fraction of each galaxy was studied. The Leo spheroidals had sufficient numbers of stars for accurate star formation history measurements; the others produced only rough star formation histories. The consistent feature of the star formation histories is that ancient ($> 8$ Gyr) star formation was detected in all eight CMDs at the $1 \\sigma$ level. After the ancient burst, some (Ursa Minor, Draco, and Sculptor) show no evidence of young star formation. Leo II shows star formation covering about half its lifetime, while Carina and Sagittarius appear to have formed stars until $\\sim 2$ Gyr ago. Finally, Leo I shows a very strong young burst, with its star formation rate $2-3$ Gyr ago nearly four times its lifetime average. Results for the galaxies are summarized in Table \\ref{tab_summary}." + }, + "0112/astro-ph0112107_arXiv.txt": { + "abstract": "The surface electric characteristics of bare strange stars are investigated with the inclusion of boundary effects. The thickness of the electron layer where pairs can be spontaneously created is calculated as a function of the bag parameters. We find that previous estimates are representative for bag parameters within a rather wide range, and therefore our results favor the thermal radiation mechanism of bare strange stars advanced by V. V. Usov. ", + "introduction": "It is of great importance to identify strange stars; a new window of distinguishing neutron stars and bare strange stars (BSSs) has been proposed recently according to their sharp differences of surface conditions (Xu, Zhang, Qiao 2001; Usov 2001a). It is therefore essential to study the surface properties of BSSs in details, e.g., the degree of the thermal luminosity of a hot BSS. The surface electric field should be very strong ($\\sim 10^{17}$ V/cm) near the bare quark surface of a strange star because of the mass difference of the strange quark and the up (or down) quark (Alcock et al. 1986), which could play an important role in producing the thermal emission of BSSs by the Usov mechanism (Usov 1998, Usov 2001b), because the strange quark matter is a poor radiator of thermal photons at frequencies less than its plasma frequency ($\\sim 20$ MeV) (Alcock et al. 1986). The basic idea of the Usov mechanism is that $e^{\\pm}$ pairs are created rapidly in a few empty quantum states with energy $\\epsilon<\\epsilon_{\\rm F}-2mc^2$ ($\\epsilon_{\\rm F}$ is the Fermi energy, $m$ the electron mass) due to the very strong electric field in an electron layer (with a height of $\\sim 500$ fm above quark surface); the pairs subsequently annihilate into photons which are then thermalized in the electron layer\\footnote{ % This layer may be optically thick for BSSs with temperature $T\\ga 10 ^9$ K.} around a BSS. This radiative mechanism has been applied tentatively to soft $\\gamma$ ray repeaters (Usov 2001c; Usov 2001d) recently. In addition, the strong electric field plays an essential role in forming a possible crust around a strange star, which has been investigated extensively by many authors (e.g., Martemyanov 1992; Kettner et al. 1995; Huang \\& Lu 1997; Phukon 2000; see Zdunik, Haensel \\& Gourgoulhon 2001 for the recent development with the inclusion of rotating and general-relativistic effects). Also it should be noted that this electric field has some important implications on pulsar radio emission mechanisms (Xu \\& Qiao 1998; Xu, Qiao, Zhang 1999; Xu, Zhang, Qiao 2001). In fact the Usov mechanism of pair production depends on many parameters; it is therefore imperative to study the dependence of the process on these parameters. With some typical parameters chosen by Usov in his calculations, the resultant thickness of the electron layer with electric field $E\\ga 1.3\\times 10^{16}$ V/cm (the critical field necessary for pair production), $\\Delta r_{\\rm E}$, is $\\sim 500$ fm. However the proper determination of the thickness $\\Delta r_{\\rm E}$ should be done with the dynamical theory of the strange quark matter. Because of the untractable nature of the quantum chromodynamics, some phenomenological models, i.e., the MIT bag model (e.g., Jensen \\& Madsen 1996), the quark mass-density-dependent model (e.g., Lugones \\& Benvenuto 1995), and the quark potential model (e.g., Dey et al. 1998), have been applied to the descriptions of the strange quark matter . In the bag model, $\\Delta r_{\\rm E}$ is a function of $\\alpha_{\\rm c}$ (the coupling constant for strong interactions), $m_{\\rm s}$ (the strange quark mass), and $B$ (the bag constant). Also it should be noticed that the quark number densities, $n_{\\rm i}$ (i = u, d, s for up, down and strange quarks, respectively), are assumed to be uniform below the quark surface, and therefore the quark charge density $(2n_{\\rm u}-n_{\\rm d}-n_{\\rm s})/3=V_{\\rm q}^3/(3\\pi^2)$ is constant near the surface. The potential $V_{\\rm q}$ is usually chosen as 20 MeV for typical cases. However the quark number densities should not be uniform near the surface if boundary effects are included, since charge neutrality is broken there. In this paper we improve the calculation of the electric field in the vicinity of a BSS surface, using the popular bag model and the Thomas-Fermi model. We investigate the electric characteristics of BSSs for different values of $\\alpha_{\\rm c}$. Initially we use typical parameters for $m_{\\rm s}=200$ MeV and $B^{1/4}\\sim 145$ MeV. Then the thickness, $\\Delta r_{\\rm E}$, of the electron layer where pairs can be created, is computed as functions of $B$, $m_{\\rm s}$, and $\\alpha_{\\rm c}$. It is found that Usov's estimates are representative for bag-model parameters within a rather wide range. Boundary effect are also considered in this calculation. ", + "conclusions": "We have improved the calculation of the electric characteristics of bare strange stars with the inclusion of boundary effects (i.e., the effects of non-local neutrality near and below the quark surface). From our calculation, we find that the Usov mechanism can work for bag parameters within a rather wide range. As shown in Table 1, both $V_{\\rm c}$ and $V_0$, as well as their very small difference, decrease as $\\alpha_{\\rm c}$ increases. Our results on the electric potential $V(z)$ for $z<0$ is quit different from the previous calculation given by Alcock et al.(1986) where the boundary effects were not included (see Fig. 1). It is shown that the strong electric field resides only about $\\sim 10$ fm below the quark surface (see Fig. 3), rather than $\\sim 10^2$ fm obtained by Alcock et al. (1986). We can proof that $E$ is almost an exponential function of $z$ below the quark surface. Denote the right hand side of the Eq.(\\ref{a6}) as $f(V)$. As $z\\rightarrow-\\infty$, namely $V \\rightarrow V_0$, we have $f(V)\\rightarrow0$. Approximating $f(V)$ as $f'(V_0)(V-V_0)$, we can obtain $E=-\\sqrt{f'(V_0)}\\exp[\\sqrt{f'(V_0)}]z$ and $V=V_0-\\exp[\\sqrt{f'(V_0)}z]$ for $z < 0$. \\vspace{2mm} \\parindent=0pt {\\it Acknowledgments.}~~ This work is supported by National Nature Sciences Foundation of China (10173002) and the Special Funds for Major State Basic Research Projects of China (G2000077602). The authors sincerely thank Dr. Shuangnan Zhang for his comments and the improvement of the language." + }, + "0112/hep-ph0112118_arXiv.txt": { + "abstract": " ", + "introduction": "The mixing between photons and axions, in an external magnetic field, is a well studied mechanism~\\cite{raffelt88,sikivie83,sikivie84,morris} as is its analog with gravitons~\\cite{gertsenshtein62,zeldovich83}. It is experimentally used, since the pioneer work by Sikivie~\\cite{sikivie83}, to constrain the axion parameters~\\cite{bibber89,hagmann98,kim98,sikivie97,gg} (see also e.g. Ref. \\cite{raffelt98} for an up to date review on such experiments). The astrophysical and cosmological implications of this mechanism have also been studied~\\cite{raffelt98}, and it was recently advocated to be a possible explanation for the observed dimming of distant type Ia supernovae~\\cite{sn2,sn1} by Cs\\'aki {\\em et al.}~\\cite{Csaki:2001yk} (see also \\cite{Erlich:2001iq} for the analysis of supernovae data with the model of~\\cite{Csaki:2001yk}). The underlying idea is simply that the luminosity of distant supernovae can be diminished due to the decay of photons into very light pseudoscalar particles induced by intergalactic magnetic fields over cosmological distances, and thus that the luminosity distance-redshift relationship can mimic the one of a universe with a non zero cosmological constant without need for a cosmological constant. The pseudoscalar particles must have an electromagnetic coupling similar to axions, and a specific and very small mass ($m\\sim 10^{-16}$ eV) to avoid affecting the cosmic microwave background anisotropy beyond its observed value. Another aspect of photon-pseudoscalar conversion in intergalactic magnetic fields is the change of the polarization properties of distant sources~\\cite{diego} such as supernovae. The implications on the cosmic microwave background of the similar effect involving photon-graviton oscillation has been considered in empty space~\\cite{magueijo94,chen95} and it was shown that it becomes negligible for standard cosmological magnetic fields~\\cite{cillis96} once the contribution of the intergalactic plasma is properly taken into account (the case of axions is also considered in~\\cite{du1}). The effect of the inhomogeneities of the electron density upon the coherence of the oscillations was also considered by Carlson {\\em et al.}~\\cite{carl}. The effect of this plasma was not addressed in~\\cite{Csaki:2001yk}, and it is our purpose to discuss it in this work. The paper is organized as follows: we first remind standard results on photon-pseudoscalar oscillations in order to introduce our conventions (see~\\cite{raffelt88}, or~\\cite{du1} where contributions from Kaluza-Klein modes were also included). We then discuss specifically the effect of the plasma for the parameters relevant for type Ia supernovae. ", + "conclusions": "We have shown that one can not ignore the effect of the intergalactic plasma to derive how the luminosity of distant sources, such as supernovae, is affected by a mixing with a hypothetical pseudoscalar particle. In most of the parameter space, this effect either renders the oscillation frequency-dependent or lowers too much the oscillation probability. There is a slight hope to accommodate the mechanism of \\cite{Csaki:2001yk} if the IGM has very specific statistical properties." + }, + "0112/astro-ph0112195_arXiv.txt": { + "abstract": "In astrophysical situations, e.g. in the interstellar medium (ISM), neutrals can provide viscous damping on scales much larger than the magnetic diffusion scale. Through numerical simulations, we have found that the magnetic field can have a rich structure below the dissipation cutoff scale. This implies that magnetic fields in the ISM can have structures on scales much smaller than parsec scales. Our results show that the magnetic energy contained in a wavenumber band is independent of the wavenumber and magnetic structures are intermittent and extremely anisotropic. We discuss the relation between our results and the formation of the tiny-scale atomic structure (TSAS). ", + "introduction": "In the interstellar medium (ISM), flows are complicated and dynamic. Observations suggest that the ISM is in a turbulent state (Larson 1981; Myers 1983; Scalo 1984; Armstrong, Rickett \\& Spangler 1995; Stanimirovic \\& Lazarian 2001). Hydrodynamic turbulence can be described by so-called energy cascade model, in which energy injected at a scale $L$ cascades down to progressively smaller and smaller scales. Ultimately, the energy will reach the molecular {\\it dissipation scale} $l_d$ and the energy will be lost there. The scales between $L$ and $l_d$ constitute the {\\it inertial range}. In hydrodynamic turbulence the dissipation scale is the minimal scale for motion. If we plot kinetic energy spectrum $E_v(k)$, the kinetic energy contained in a wavenumber band, we will see that a) the spectrum peaks at the wavenumber corresponding to the energy injection scale ($k_L \\sim 1/L$); b) it follows a power law (e.g. $E_v(k) \\propto k^{-5/3}$ in Kolmogorov theory) in the inertial range; c) it drops rapidly after $k_d \\sim 1/l_d$, which depends on viscosity $\\nu$: $k_d \\sim (VL/\\nu)^{3/4}k_L$, where $V$ is the rms velocity at the energy injection scale $L$. Magnetohydrodynamic (MHD) turbulence has two energy loss scales - a viscous damping scale set by the viscosity $\\nu$ and magnetic diffusion scale set by the ohmic resistivity $\\eta$. When $\\nu\\gg\\eta$, the viscous damping scale is much larger than the magnetic diffusion scale. Although MHD turbulence is different from its hydrodynamic counterpart in many ways, the energy cascade model is still valid (e.g. Goldreich \\& Sridhar 1995). Therefore, in this case, kinetic energy is damped before the cascading energy reaches the magnetic diffusion scale. In this paper we study MHD turbulence when the mean field $B_0$ is at least as large as the fluctuating field $b$. This is the opposite of the dynamo regime with $B_0\\ll b$ discussed in Kulsrud \\& Anderson (1992) and tested in Maron \\& Cowley (2001). We shall show that the wide-spread assumption that magnetic structures do not exist below the viscous cutoff is wrong. In astrophysics, the viscosity caused by neutrals damps turbulence. In the ISM, this viscous cutoff occurs at $\\sim$~$pc$ scales, which is much larger than the magnetic diffusion scale. To model this, we use a large physical viscosity and very small magnetic diffusivity. In this letter we numerically demonstrate the existence of a power-law magnetic energy spectrum below the viscous damping scale. We use an incompressible MHD code. ", + "conclusions": "We have considered MHD turbulence with $\\nu \\gg \\eta$, which implies that the viscous cutoff occurs at a scale much larger than the magnetic diffusion scale. It has been believed that the damping of the fluid motion is accompanied by the suppression of magnetic structures below the viscous cutoff scale. To the contrary, we have found that the magnetic field can have a rich structure below the viscous cutoff scale. Judging from our (limited) numerical results, we conclude that magnetic field perturbations have a similar power distribution as a passive scalar, despite the obvious importance of magnetic forces. In particular, the spectrum follow a $k^{-1}$ law. The parallel wavenumber, which is an indicator of the degree of anisotropy, is almost constant. Consequently, the magnetic field has an extreme form of anisotropy. In summary, we have found that \\begin{enumerate} \\item $E_b(k)\\propto k^{-1}$, \\item $E_v(k)\\propto k^{-4}$, \\item $k_{\\|} \\approx \\mbox{constant.}$ \\end{enumerate} We discussed the possibility that this small scale magnetic field is the cause of the tiny-scale atomic structure. The small scale magnetic structure will affect many astrophysical processes (e.g. cosmic-ray transport, reconnection, etc) that depend on the statistical properties of MHD turbulence." + }, + "0112/astro-ph0112476_arXiv.txt": { + "abstract": "We present predictions of the correlation between the \\lya forest absorption in quasar spectra and the mass within $\\sim 5 \\hmpc$ (comoving) of the line of sight, using fully hydrodynamic and hydro-PM numerical simulations of the cold dark matter model supported by present observations. The observed correlation based on galaxies and the \\lya forest can be directly compared to our theoretical results, assuming that galaxies are linearly biased on large scales. Specifically, we predict the average value of the mass fluctuation, $<\\! \\dm \\! >$, conditioned to a fixed value of the \\lya forest transmitted flux $\\dF$, after they have been smoothed over a $10 \\hmpc$ cube and line of sight interval, respectively. We find that $<\\! \\dm \\! >/\\sigma_m$ as a function of $\\dF/\\sigma_F$ has a slope of $0.6$ at this smoothing scale, where $\\sigma_m$ and $\\sigma_F$ are the rms dispersions (this slope should decrease with the smoothing scale). We show that this value is largely insensitive to the cosmological model and other \\lya forest parameters. Comparison of our predictions to observations should provide a fundamental test of our ideas on the nature of the \\lya forest and the distribution of galaxies, and can yield a measurement of the bias factor of any type of galaxies that are observed in the vicinity of \\lya forest lines of sight. ", + "introduction": "The prevalent theory to explain the \\lya forest is that the absorption lines arise from density variations in a photoionized intergalactic medium that originate in the gravitational evolution of primordial fluctuations. Both semi-analytic models and numerical simulations (Bi 1993; Cen \\etal 1994; Zhang \\etal 1995, 1998; Hernquist \\etal 1996; Miralda-Escud\\'e \\etal 1996; Bi \\& Davidsen 1997) have shown that the predicted \\lya spectra appear remarkably similar to the observations. The good agreement of the predicted and observed flux distribution and power spectrum of the \\lya forest (Rauch \\etal 1997; Croft \\etal 1999; McDonald \\etal 2000), and the large transverse size of the absorption systems (Bechtold \\etal 1994; Dinshaw \\etal 1994, 1997; Petitjean \\etal 1998; Monier, Turnshek, \\& Hazard 1999; Dolan et al. 2000; L\\'opez, Hagen, \\& Reimers 2000) are the basic tests that have so far been done and have supported the theory. In addition, this \\lya forest theory is derived from the general Cold Dark Matter (hereafter, CDM) model with parameters that are well constrained from several other observations (e.g., Primack 2000). Another important test of our ideas of the \\lya forest can be done by observing the correlation of the transmitted flux in a spectrum with galaxies. Some of these observations have already been done for small scales and high column density absorbers, which have shown the expected strong correlation between galaxies and gas halos (e.g., Bergeron \\& Boiss\\'e 1991; Steidel, Dickinson, \\& Persson 1994; Lanzetta \\etal 1995; Chen \\etal 2001). Penton, Stocke, \\& Shull (2001) have probed this correlation at low redshift, and have found that the weak absorption lines are often found in low-density regions of the galaxy distribution. Recently, Adelberger \\etal (2001) have carried out the first analysis of this correlation on large scales (several comoving Mpc) and at high redshift (using galaxies detected with the Lyman break technique), using the transmitted flux as the quantity to correlate with the mean number of galaxies in a specified region. Inspired by the observational results of Adelberger \\etal (2001), this paper presents detailed theoretical predictions for statistical functions similar to the one introduced by those authors. Namely, we analyze the mean value of the mass given an observed value of the transmitted flux, and the mean value of the transmitted flux for a fixed value of the mass, after both quantities have been smoothed over a certain region. The main difference between the functions we analyze and the function shown by Adelberger \\etal (2001) is that the observed objects are of course galaxies, and our simulations predict only the distribution of the mass. However, the predictions of the simulations can still be compared to the galaxy observations to determine the relation between galaxies and mass, which is especially straightforward if linear bias is a sufficient description of the distribution of galaxies relative to the mass on the large scales being probed. We examine the dependence of these statistical functions on the various parameters affecting the \\lya forest in \\S 3. ", + "conclusions": "We have presented predictions for the expected correlation of the mass and the \\lya forest transmitted flux, smoothed over a cube size of $\\sim 10 \\hmpc$. Our most basic result is that the relation between $<\\! \\dtm | \\dtF\\! >$ and $\\dtF$ is linear over most of the range of $\\dtF$ (excluding rare, high $\\dtF$ values), with a slope of 0.6 on this smoothing scale. We have shown that this correlation is not sensitive to the temperature-density relation of the gas. There is a dependence on the mean transmitted flux and the amplitude of the mass power spectrum, but these quantities are already measured to reasonable accuracy (Croft \\etal 1999, McDonald \\etal 2000). We have also shown that the predictions do not suffer from uncertainties due to the resolution of the simulations. Calculations with even larger boxes than used here will be desirable, because the uncertainties in the predictions due to the variance in the simulations and the suppression of the large-scale power may still be significant; however, this uncertainty is mostly isolated in the most rare, high density regions. A similar relation holds between $<\\! \\dtF | \\dtm\\! >$ and $\\dtm$. We notice, however, that the observational determination of this other function will be affected by galaxy shot noise. The number of galaxies observed in a certain cube in redshift space containing a fixed mass is subject to shot noise, and this inevitably introduces a smoothing of the function $<\\! \\dtF | \\dtm\\! >$, which must be taken into account before any comparisons to our theoretical results are made (in contrast, galaxy shot noise does not change the average $<\\! \\dF | \\dm \\! >$, but it does alter the rms dispersion $\\sigma_F$ needed to compute $\\dtF$). Measuring the effects of galaxy shot noise can also teach us useful information about how galaxies form, because galaxy shot noise does not generally need to be strictly described by Poisson statistics. For example, for a fixed mass contained within a cube, once a galaxy is found in the cube the probability to find others may be lower because some mass has already been used up by that galaxy. Comparing these theoretical predictions with observations allows for several tests of the basic theory of the \\lya forest, and can reveal new information on the spatial distribution of the galaxies. The functions $<\\! \\dtm | \\dtF \\! >$ and $<\\! \\dtF | \\dtm \\! >$, which should not be affected by any linear galaxy bias, provide a powerful test of the assumption that the basic framework assumed here for the nature of the \\lya forest is correct, and that galaxies trace the mass on large scales apart from linear bias. If the slope of $<\\! \\dtm | \\dtF \\! >$ as a function of $\\dtF$ were found to be larger than predicted, this would imply that the \\lya forest is much more closely associated with galaxies than is expected from their common correlation with the mass distribution. If the observed slope were smaller than predicted, it would indicate that the \\lya absorbing gas and the galaxies tend to avoid each other for some reason. If the correlation of $\\dtm$ and $\\dtF$ is as expected, this will imply a strong confirmation of the basic model we have for the \\lya forest, and will justify using the \\lya forest as a predictor of the mass fluctuations. Comparing the predicted function $<\\! \\dm | \\dF \\! >$ with the observed $<\\! \\dg | \\dF \\! >$ will then yield the bias factor of any type of galaxies for which these observations can be made. The correlation of galaxies and the \\lya forest can be measured as a function of scale. Our predicted value of 0.6 for the slope of $< \\! \\dtm | \\dtF \\! > $ as a function of $\\dtF$ should decrease with the smoothing scale, in a way that reflects the shape of the mass autocorrelation function. This will allow a precise test of the idea that the large-scale distributions of different types of galaxies differ only in a constant linear bias relative to the mass fluctuations." + }, + "0112/math-ph0112040_arXiv.txt": { + "abstract": "I investigate the properties of forces on bodies in theories governed by the generalized Poisson equation $\\div[\\m(\\abgf/\\e)\\gf]\\propto G\\r$, for the potential $\\f$ produced by a distribution of sources $\\r$. This equation describes, {\\it inter alia}, media with a response coefficient, $\\m$, that depends on the field strength, such as in nonlinear, dielectric, or diamagnetic, media; nonlinear transport problems with field-strength dependent conductivity or diffusion coefficient; nonlinear electrostatics, as in the Born-Infeld theory; certain stationary potential flows in compressible fluids, in which case the forces act on sources or obstacles in the flow. The expressions for the force on a point charge is derived exactly for the limits of very low and very high charge. The force on an arbitrary body in an external field of asymptotically constant gradient, $-\\vg_0$, is shown to be $\\vF=Q\\vg_0$, where $Q$ is the total effective charge of the body. The corollary $Q=0 \\Rightarrow \\vF=0$ is a generalization of d'Aembert's paradox. I show that for $G>0$ (as in Newtonian gravity) two point charges of the same (opposite) sign still attract (repel). The opposite is true for $G<0$. I discuss the generalization of this to extended bodies, and derive virial relations. \\vskip 9pt PACS numbers: ", + "introduction": "\\setcounter{equation}{0} \\par The Poisson equation, which governs so many physical processes, has the nonlinear generalization \\beq \\div[\\m(\\abgf/\\e)\\gf]=\\ad G\\r, \\label{i} \\eeq by which the source distribution $\\r(\\vr)$, in $D$-dimensional Euclidean space, gives rise to a potential field $\\f$. Here, $\\e$ is a constant with the dimensions of $\\gf$, $\\ad=2(\\pi)\\^{D/2}/\\Gamma(D/2)$ is the $D$-dimensional complete solid angle, introduced here for convenience, and $G$ is a coupling constant. As I will show, for $ G>0$, a point, test charge is attracted to a (finite) point charge of the same sign (as in gravity), while for $ G<0$ it is repelled. \\par Equation(\\ref{i}) describes a variety of physical problems; some examples are: \\par (i) Nonlinear dielectric, and diamagnetic, media; $\\m$ is then the dielectric or diamagnetic coefficient, which depends on the field strength (here $G<0$). \\par (ii) Problems of nonlinear electric-current flows in systems with field-dependent conductivity (nonlinear current-voltage relation), and nonlinear diffusion problems; $\\m(\\ma)$ is the transport coefficient. \\par (iii) Stationary, subsonic, potential-flow problems of non-viscous fluids with a barotropic equation of state $p=p(\\varrho)$ ($p$ is the pressure, $\\varrho$ the density). The stationary Euler equation is integrated into Bernoulli's equation $f(\\varrho)=-{1\\over 2}u^2 + const.$, where $f'(\\varrho)=\\varrho^{-1}p'(\\varrho)=c^2(\\varrho)/\\varrho$, with $c$ the speed of sound; $f$ is thus increasing with $\\varrho$, and $\\varrho$ is a function of $\\abs{\\vu}=\\abs{\\gf}$. The stationary continuity equation then gives $ \\div[\\varrho(\\abgf)\\gf]=s(\\vr)$, with $s$ the source density (see e.g. \\cite{gt} for the ideal-gas case). This is eq.(\\ref{i}) with $\\m=\\varrho$, and $G=\\ad^{-1}>0$. For example, in a fluid with an equation of state of the form $p=a \\varrho^{\\c}$ ($a>0,~\\c\\ge 1$), we have $\\varrho(u)=\\varrho(0) [1-(u/u_0)^2]^{1/(\\c-1)}$, with $u_0^2\\equiv 2c_0^2/(\\c-1)$, and $c_0$ is the speed of sound at $u=0$. Subsonicity requires $(u/u_0)^2<(\\c-1)/(\\c+1)$. \\par If we, formally, consider a stationary flow problem in a medium with negative compressibility, $c^2<0$, ellipticity is maintained for any value of $\\gf$. For example, for a medium with a constant $c^2<0$, $\\varrho(u)=\\varrho(0)exp(u^2/2\\abs{c}^2)$. \\par (iv) Nonlinear (vacuum) electrostatics as formulated e.g. in the Born-Infeld nonlinear electromagnetism, which also appears in effective Lagrangians resulting from string theory (see review and references in \\cite{gibra}\\cite{gib}). In the original, electrostatic Born-Infeld theory $\\m(\\ma)\\propto (1-\\ma^2)^{-1/2}$, and $G<0$. \\par (v) A formulation of an alternative nonrelativistic gravity to replace the dark-matter hypothesis in galactic systems \\cite{bm}. Here $\\m(\\ma)\\approx \\ma$ for $\\ma\\ll 1$, and $\\m\\approx 1$ for $\\ma\\gg 1$ ($G>0$). \\par (vi) Equation(\\ref{i}) was used in \\cite{adl} as an effective-action approximation to Abelianized QCD. \\par (vii) Area (volume) minimization problems, such as the determination of the shape of a soap film with a dictated boundary (see e.g. \\cite{gt}): If $x\\_{D+1}=\\f(x_1,...,x\\_{D})$ describes a $D$-dimensional hypersurface embedded in $D+1$ dimensional Euclidean space with Cartesian coordinates $x_1,...,x\\_{D+1}$, the volume element on the surface is $dv=[1+(\\gf)^2]^{1/2}~d^Dr$. Then, eq.(\\ref{i}) describes the problem of the minimization of the volume of the surface. The sources may describe a force density on the hypersurface acting in the direction $x\\_{D+1}$. In this problem $\\m(\\ma)\\propto (1+\\ma^2)^{-1/2}$, and $G>0$. Born-Infeld electrostatics is the same as the area-extremization problem for a time surface embedded in Minkowski space-time. \\par Much has been said in the mathematical literature on properties of the potentials that solve eq.(\\ref{i}) (see e.g. \\cite{gt}). But, to my knowledge, very little has been said about forces on bodies in such theories. The forces can be written as certain integrals of $\\f$, and are of obvious relevance in the physics context. \\par Many of the familiar and intuitive properties of the linear theory are lost in the nonlinear case because the potential, and forces, is not the sum of the contributions of the sub-systems. For example, the force on a point charge is no more proportional to the charge, does not reverse direction when the charge of the body reverses sign, etc.. And, Earnshaw's theorem \\cite{earnshaw} no longer holds. Under some circumstances it is possible to suspend stably static charged bodies in a static field. This last aspect is treated in detail in \\cite{sus}. \\par The choice $\\m(\\ma)=\\ma^{D-2}$ is special in that the theory is then conformally invariant, and lands itself to many analytical developments. This case is described in detail in \\cite{conf}. \\par After discussing some general aspects of eq.(\\ref{i}) in section II, I take up the main subject concerning the properties of forces on bodies in the nonlinear theories: general properties in sections III, and forces on point charges in section IV. I conclude in section V with some examples of applications. ", + "conclusions": "" + }, + "0112/astro-ph0112530_arXiv.txt": { + "abstract": "We present a comparison of X-ray and optical luminosities and luminosity functions of cluster candidates from a joint optical/X-ray survey, the ROSAT Optical X-ray Survey (ROXS). Completely independent X-ray and optical catalogs of 23 ROSAT fields (4.8 square degrees) were created by a matched-filter optical algorithm and by a wavelet technique in the X-ray. We directly compare the results of the optical and X-ray selection techniques. The matched-filter technique detected 74\\% (26 out of 35) of the most reliable cluster candidates in the X-ray-selected sample; the remainder could be either constellations of X-ray point sources or $z>1$ clusters. The matched-filter technique identified approximately 3 times the number of candidates (152 candidates) found in the X-ray survey of the same sky (57 candidates). While the estimated optical and X-ray luminosities of clusters of galaxies are correlated, the intrinsic scatter in this relationship is very large. We can reproduce the number and distribution of optical clusters with a model defined by the X-ray luminosity function and by an $L_x-\\Lambda_{cl}$ relation if $H_0=75$~km~s$^{-1}$~Mpc$^{-1}$ and if the $L_x-\\Lambda_{cl}$ relation is steeper than the expected $L_x \\propto \\Lambda_{cl}^2$. On statistical grounds, a bimodal distribution of X-ray luminous and X-ray faint clusters is unnecessary to explain our observations. Followup work is required to confirm whether the clusters without bright X-ray counterparts are simply X-ray faint for their optical luminosity because of their low mass or youth, or a distinct population of clusters which do not, for some reason, have dense intracluster media. We suspect that these optical clusters are low-mass systems, with correspondingly low X-ray temperatures and luminosities, or that they are not yet completely virialized systems. ", + "introduction": "X-ray and optical techniques for detecting clusters of galaxies each have their merits. Optical techniques have been in use for over four decades, and images of the optical sky are relatively inexpensive to obtain. New optical methods such as the matched-filter method (Postman et al. 1996, P96 hereafter) have allowed for automatic, uniform detection of optical overdensities of galaxies with magnitude distributions consistent with those of a typical cluster of galaxies. In the matched-filter method, fitting galaxy luminosity functions in addition to seeking overdensities of galaxies minimizes the projection effects that plagued earlier cluster cataloguing efforts. X-ray selection has the advantage of directly revealing the hot intracluster medium confined by the deep gravitational potential of the cluster. Since this emission is proportional to the gas density squared, and the X-ray sky is sparsely populated compared to the optical, X-ray detections have higher contrast and less contamination from physically unrelated systems. Early X-ray selection methods optimal for point sources (Gioia et al. 1990b) were biased somewhat towards selecting clusters with high central surface brightneses, but there are now several algorithms optimized for detecting extended sources, including wavelets and Voronoi-Tesselation Percolation methods (Rosati et al. 1995; Scharf et al. 1997). However, X-ray and optical studies of cluster evolution have progressed along separate paths. A decade ago, optical and X-ray surveys disagreed in their assessment of the amount of evolution clusters have experienced. Optical surveys indicated very little evolution since $z\\sim0.5-1$ (Gunn, Hoessel \\& Oke 1986) while X-ray studies suggested modest (Gioia et al. 1990a) to strong evolution (Edge et al. 1990, later retracted in Ebeling et al. 1997). The most recent X-ray samples of clusters over a range of redshifts out to $z\\sim0.8-1.2$ suggest that the X-ray luminosity function for moderate luminosity clusters has in fact not evolved significantly since $z\\sim0.8$ (Borgani et al. 1999; Nichol et al. 1999; Rosati et al. 1998; Jones et al. 1998), while the {\\em most luminous} (and presumably most massive) systems, such as those contained in the EMSS, may have evolved somewhat (Gioia et al. 1990a; Henry et al. 1992; Nichol et al. 1997, Vikhlinin et al. 1998, 2000; Rosati et al. 2000; Gioia et al. 2001), but the community is not unanimous on this result (e.g. Jones et al. 1998). In contrast, recent optical surveys for distant clusters continue to find very little evidence for evolution (Couch et al. 1991; P96). In an effort to establish the common ground between X-ray and optical studies of clusters and cluster evolution, we have undertaken a joint X-ray/optical survey to detect and study clusters of galaxies, called the ROSAT Optical X-ray Survey, or ROXS. In contrast to other ROSAT PSPC serendipitous surveys (e.g. Rosati et al. 1995, Jones et al. 1998, Romer et al. 2000, Vikhlinin et al. 1998), the ROXS team optically imaged the entire central 30' by 30' of each field in our survey. The X-ray selection and optical selection of cluster candidates were then determined independently. In this current work, we report results on the relation between the X-ray luminosity ($L_x$) and $\\Lambda_{cl}$, a measure of the cluster's optical luminosity. The ROXS has already been used to find a potential intergalactic X-ray filament (Scharf et al. 2000.) Paper II (Donahue et al., 2001) contains the catalogues, data reduction and observational details. For this paper we scale $H_0=75~h_{75}$ km/s/Mpc. We assume $q_0=0.5$ to ease comparison with earlier results. ", + "conclusions": "We have directly compared the results of cluster hunting with two competing methods: optical selection in the I-band, using the matched filter technique, and X-ray selection of extended sources. We have found that both methods reliably detect most of the richest, and presumably most massive systems, but the optical matched filter technique, because of the scatter and the steepness of the relation between X-ray and optical luminosities, produces more cluster candidates at our sensitivities. We present the first $\\Lambda_{cl}$ function for optically-selected clusters of galaxies. The $\\Lambda_{cl}$ function is consistent with the X-ray luminosity function for clusters of galaxies if the intrinsic, global $\\Lambda_{cl}-L_x$ relation is consistent with that of the cross-identified candidates. In particular, we can explain our observations without appealing to an X-ray faint population of massive clusters of galaxies. Predictions for an observational test using deep XMM-Newton/EPIC exposures based on our result are: (1) X-ray observations of the undetected, high-$\\Lambda_{cl}$ systems should detect X-rays from most of the systems with observations of only moderately increased sensitivity over the ROSAT observations, and the detected ICM will not be hot ($T_x\\lesssim4$ keV). (2) If $L_x \\propto \\Lambda_{cl}^{3-4}$, the median predicted bolometric X-ray luminosity for the entire undetected sample is $\\sim 10^{43}h_{75}^{-2}$ erg s$^{-1}$, so that deep XMM-Newton EPIC exposures should detect most of this population. On the other hand, significant numbers of non-detections with XMM will suggest that many of these clusters are spurious or are significantly less massive than their estimated richnesses may suggest or that there is an X-ray faint population at these mass scales. (3) The X-ray candidates which are not detected by the matched filter algorithm in the optical images will be revealed to be either false clusters of galaxies (constellations of AGN), moderately distant groups, or distant clusters of galaxies ($z>1-1.2$)." + }, + "0112/astro-ph0112030.txt": { + "abstract": "{ X-ray emission processes in starburst galaxies (SBGs) are assessed, with the aim of identifying and characterizing the main spectral components. Our survey of spectral properties, complemented with a model for the evolution of galactic stellar populations, leads to the prediction of a complex spectrum. Comparing the predicted spectral properties with current X-ray measurements of the nearby SBGs M82 and N253, we draw the following", + "introduction": "In starburst galaxies (SBGs) enhanced star formation activity (lasting typically $\\mincir 10^8$ yr) drives a chain of coupled stellar and interstellar (IS) phenomena that are manifested as intense far-infrared (FIR) and X-ray emission. The SBG class is a heterogeneous group of galaxies that are selected based on optical, UV and FIR properties. Historically, optically selected galaxies, like HII galaxies, were first recognized to undergo a burst of star formation (Searle et al. 1973); the term \"starburst galaxy\" was introduced by Weedman (1981; see also Balzano 1983). Subsequently, FIR-luminous galaxies were also recognized to be SBGs, a consequence of efficient heating of IS dust by the radiation from abundant massive stars (Soifer et al. 1986). Increased stellar activity leads also to a higher supernova (SN) rate, shock heating of IS gas, and a more efficient particle acceleration mechanism in SBGs compared to `normal' spirals. Interest in SBGs stems also from the realization that these resemble young galaxies in the earlier universe. Indeed, a SB phase was very common in the earlier universe, as the cosmic star formation rate (and hence the cosmic chemical enrichment) was substantially higher at epochs corresponding to $z \\magcir 1$ (with the data being interpreted with the SFR having either a peak in the redshift range $1 \\mincir z \\mincir 2$ (Madau et al. 1996), or -- due to small number statistics and dust correction effects -- a plateau out to $z \\sim 4$ (Thompson et al. 2001). So, if the main properties of SBGs in the present universe resemble those of normal galaxies during the evolutionary phase at $z > 1$, the study of local SBGs may provide insight into processes that occurred at that earlier epoch. A primary manifestation of the SB activity is X-ray emission. Given the greatly enhanced star formation rate, energetic phenomena related to the final stages of stellar evolution -- X-ray binaries, supernova remnants (SNRs), galactic winds, and Compton scattering of ambient FIR photons off SN-accelerated relativistic electrons -- clearly suggest that SBGs are typically more powerful X-ray emitters than normal galaxies. In fact, normalized to the 7000 \\AA{} flux (i.e., gauging activity by the old stellar population), in the X-ray band SBGs (as well as of other classes of active galaxies) are brighter than normal (spiral and elliptical) galaxies (Schmitt et al. 1997). The mean X-ray spectrum of SBGs is expected to reflect the diverse nature of high energy activity in SBGs. The earliest attempt to determine a mean broad-band spectrum of SBGs (based on {\\it Einstein}/IPC, {\\it HEAO1}-A2, and {\\it HEAO1}-A4 data for a sample of 51 FIR-selected putative SBGs) yielded some evidence that the (co-added) emission from SBGs was detectable beyond $\\sim$10 keV even based on limited-quality survey data, and that the mean SBG spectrum was somewhat harder (photon spectrum index $\\Gamma$ $\\sim$ 1.5) than the mean AGN spectrum (Rephaeli et al. 1991, 1995). Individual spectral studies (based on data of limited spatial resolution) have shown the 0.5-10 keV spectra of SBGs to be complex: these are best-fit by one (or more) low-temperature ($kT < 1$ keV) component(s) plus a harder component, the latter interpreted as being either thermal with $kT \\sim 5-10$ keV or nonthermal with $\\Gamma \\sim 1.5-2$ (M82, N253: Ptak et al. 1997, Cappi et al. 1999; M83: Okada et al. 1997; N2146: Della Ceca et al. 1999; N3256: Moran et al. 1999; N3310, N3690: Zezas et al. 1998). A purely thermal hard component would imply low chemical abundance ($Z \\sim 0.3 \\, Z_\\odot$), whereas comparable contributions of thermal and nonthermal emissions would imply $Z \\sim Z_\\odot$ % \\footnote{Solar abundances can also be recovered using different spectral models (Weaver et al. 2000). More generally, however, it should be emphasized that abundance determinations are difficult due to the uncertainties in the Fe-L atomic physics, because Fe-L lines couple with O and Ne-K lines upon which abundance determinations rely strongly (e.g., Matsushita et al. 2000).}. % The issue of whether the hard component is actually thermal with inferred high temperatures and strongly subsolar abundance, and whether it originates from genuinely hot diffuse gas or from unresolved point sources, is not settled yet (e.g.: Weaver et al. 2000, Dahlem et al. 2000). Furthermore, the spectrum of at least some SBGs may also include a substantial contribution from a compact nuclear source: see, e.g., the apparent temporal variability of the 2-10 keV flux found in {\\it ASCA} and {\\it RXTE} measurements of M82 (Tsuru et al. 1997, Ptak \\& Griffiths 1999, Rephaeli \\& Gruber 2001). More recently, for N253, {\\it Chandra} data have resolved the regions where the soft thermal X-rays originate (Strickland et al. 2000), and {\\it XMM} has separated the extended and point-like emission components in the disk and the nuclear region: the unresolved (diffuse?) emission is spatially and spectrally complex, with two warm plasma components in the disk and three (warm and hot) in the nucleus (Pietsch et al. 2001). For both N253 and M82, {\\it Chandra} data have shown that the $\\sim$ 2-10 keV flux is dominated by point-source emission (Strickland et al. 2000, Griffiths et al. 2000). The main goal of this paper is an attempt to quantify the X-ray characteristics of the (stellar and non-stellar driven) modes of activities in SBGs, in order to identify basic spectral features that will help elucidate the nature of these galaxies. Starting from a realistic stellar population model for X-ray binaries and SNRs in our Galaxy, we account for all the viable stellar and gaseous X-ray emitting processes in a galactic environment, and describe the respective spectra in Section 2. In section 3 we construct a template for the composite X-ray spectrum of a galaxy; this is then compared (in Section 4) with measurements of the nearby SBGs M82 and N253. We conclude with a summary of our main findings (Section 5). ", + "conclusions": "" + }, + "0112/astro-ph0112060_arXiv.txt": { + "abstract": "{ Recent far-infrared and submillimetre waveband observations revealed large number of Ultraluminous Infrared Galaxies (ULIGs) with infrared luminosities $> 10^{12}\\,L_{\\odot}$. These sources are proposed to lie at redshifts above one, and in normally interacting systems with very dusty environments. We discussed in a previous paper that a population with a fast evolving infrared burst phase triggered by gas-rich mergers at $z\\sim 1$ interpreted successfully the steep slope of faint IRAS $60\\,\\mu m$ source counts within the flux range of $100\\,mJy \\sim 1\\,Jy$, still leaving the infrared background level at this wavelength compatible with the upper limit from recent high energy TeV $\\gamma$ ray detection of Mrk501. To extend the model to mid and far infrared wavelengths, we adopt a reasonable template spectral energy distribution as typical for nearby infrared bright starburst galaxies ($L_{ir} <= 10^{12}\\,L_{\\odot}$), such as Arp220. We construct the SED for the dusty starburst mergers at $z\\sim 1$ by a simple dust extinction law and a thermal continuum assumption for the far-infrared emission. Since the radiation process at mid-infrared for these starburst merging systems is still uncertain, we assume it is similar to the MIR continuum of Arp220, but modify it by the observed flux correlation of ULIGs from IRAS and ISOCAM deep surveys. We show in this paper that the strong evolution of the European Large Area ISO Survey (ELAIS) at $90\\,\\mu m$, ISO $170\\,\\mu m$ and the Submillimeter deep survey at $850\\,\\mu m$ could be sufficiently accounted for by such an evolutionary scenario, especially the hump of the ISOCAM $15\\,\\mu m$ source count around $0.4\\,mJy$. From current best fit results, we find that the dust temperature of those extremely bright starburst merging system at $z\\sim 1$ would be higher than that of Arp220 for a reconciliation of the multi-wavelength infrared deep surveys. We thus propose that the infrared burst phase of dusty starburst galaxies or AGNs from gas-rich mergers at $z\\sim 1$ could contribute significantly to the strong evolution of the IRAS $60\\,\\mu m$, the ISO $15\\,\\mu m$, $90\\,\\mu m$, $170\\,\\mu m$, as well as the SCUBA $850\\,\\mu m$ number counts, while being compatible with the current observational limits of cosmic infrared background and the redshift distributions. The major difference of our current model prediction is that we see a fast convergence of the differential number counts at $60\\,\\mu m$ below $50\\,mJy$, which is about a factor of two brighter than other model predictions. Future infrared satellites like Astro-F or SIRTF would give strong constraints to the models. ", + "introduction": "There has been much progress in the study of extragalactic evolution since far-infrared and submillimeter deep surveys detected a significant population of Ultraluminous Infrared Galaxies (ULIGs: $L_{ir} > 10^{12}L_{\\odot}$) at high redshift ($z \\sim 1-4$) (Hughes et al. 1998, 2000, Blain et al. 1999, Eales et al. 1999, Holland et al. 1999, Puget et al. 1999, Sanders 1999, Dole et al. 2000). The major interests of present research are the nature and evolution of these sources. Due to the lack of high resolution morphological studies, the origin of these faint SCUBA sources are still not clearly understood. Local ULIGs most likely arise from major galaxy mergers with high dust extinction for the central starburst or AGN activities (Sanders \\& Mirabel 1996, Murphy et al. 1996, Surace et al. 1998, 2000, Scoville et al. 2000). Nevertheless, more than 50$\\%$ of the high redshift ultraluminous infrared objects and the optical counterparts of ISOCAM HDF-N galaxies are suggested to show indication of galactic interactions and merger signatures (Mann et al. 1997, Smail et al. 1999, Ivison et al. 1998, 2000, Sanders 1999). Meanwhile, estimation of their spectral energy distributions suggests that these galaxies are probably the high redshift counterpart of the local ULIGs discovered by the IRAS deep survey (Barger et al. 1998, 1999, Smail et al. 1998, Trentham et al. 1999, Frayer et al. 2000). Although the exact mechanism for formation of these interesting sources are not clear, there are certain reasons to say that majority of them may be major mergers of gas-rich disks accompanied by dust-shrouded nuclear starbursts or powerful Active Galactic Nuclei (AGN). SMM J14011+0252, and ERO J164502-4626.4 (HR10) are two such candidates with their central activities heavily hidden by dust extinction, and are suggested to be consistent with the evolutionary track of mergers--starbursts/AGN, probably elliptical galaxies in formation (Graham \\& Dey 1996, Cimatti et al. 1998, 1999, Frayer et al. 1998, 1999, Dey et al. 1999, Papadopoulous et al. 2001). On the other hand, the source counts from present infrared and submilimetre surveys, such as IRAS, ISO and SCUBA all significantly exceed the non-evolving predictions. The extremely strong evolution is seen from the differential counts of the ISOCAM at $15\\,\\mu m$, with the remarkable upturn at $S_{15}< 3\\,mJy$ and a fast convergence when $S_{15}\\sim 0.3\\,mJy$. This striking feature is based on the data from several independent sky surveys (Elbaz et al. 1999, Chary \\& Elbaz 2001, Mazzei et al. 2001, Serjeant et al. 2001). Although there are many other possible evolutionary scenarios which could explain the present observations, the reason that we are encouraged to explore here a merger-driven galaxy evolution picture with binary aggregation dynamics is simply because the IRAS database and recent ISO and sub-mm deep surveys indicate that most of the luminous infrared sources are actually interacting/mergering systems. Also, the local IR luminosity function shows an excess over the Press-Schechter formula (Press \\& Schechter 1974, Lonsdale 1995, Pearson \\& Rowan-Robinson 1996, Guiderdoni et al. 1998, Roche et al. 1998, Rowan-Robinson et al. 1998, Dey et al. 1999, Sanders 1999, Dole et al. 2000, Efstathiou et al. 2000, Silk \\& Devriendt 2000, Serjeant et al. 2001, Takeuchi et al. 2001). Considering mergers as a possible formation mechanism of Ultraluminous Infrared Galaxies both at high and low redshift, as well as their significant infrared emissions, Wang (1999) and Wang \\& Biermann (2000) discussed the effects of galaxy mergers on the strong evolution of the IRAS $60\\,\\,\\mu m$ deep survey within a binary aggregation galactic evolutionary scheme. In this model, the bright tail of the infrared luminosity function is simulated in a consistent way for both the density and luminosity evolution due to the decrease of the merger fration with cosmic time and a merger-triggered infrared burst phase. They found that a luminosity-dependent infrared burst phase is crucial for the interpretation of the steep slope within a flux range of $10\\,mJy\\sim 1\\,Jy$ by the IRAS $60\\,\\,\\mu m$ deep survey. This means dusty starburst galaxies or AGNs from gas rich mergers at high redshift may experience an infrared burst phase around a transition redshift $z\\sim 1$, and fade quickly within the merger time scale of that epoch. The more massive merger systems could have such infrared emission enhanced to a higher level and decrease even faster. This kind of speculation is based on the observation that ULIGs are usually more than a factor of 20 brighter than normal starburst galaxies. Although the detailed mechanism for such enormous infrared emission is still unclear, it is believed to be related to a special stage of the merger process when the dust mass and temperature are both dramatically increased (Kleinmann \\& Keel, 1987, Taniguchi \\& Ohyama 1998). Recent numerical simulation on the evolution of dusty starburst galaxies by Bekki \\& Shioya (2001) shows that there is a very strong photometric evolution during the merger process of two gas-rich disks, and a dramatic change of the spectral energy distribution (SED) around a cosmic time scale $T \\sim 1.3\\, Gyr$, when the two disks of the merger become very close and suffer from violent relaxation and the star formation becomes maximal ($\\sim 378\\,M_{\\odot}yr^{-1}$). The infrared flux in this case could increase by one magnitude, especially for the far infrared wavelength range ($60\\,\\mu m \\,\\sim 90\\,\\mu m$) in the emitting frame. The redshift distribution of the contributing sources for the steep slope at faint IRAS $60\\,\\mu m$ counts in the model of Wang \\& Biermann (2000) shows that the infrared burst phase around $z\\sim 1$ could have comparable significance to the local IR sources. The question is then whether such an infrared burst phase, or such a population of ULIGs, could also sufficiently account for the strong evolution seen in other infrared wavelengths, especially at the ISOCAM $15\\, \\mu m$, ISOPHOT $90\\,\\mu m$, $170\\,\\mu m$ and SCUBA $850\\,\\mu m$. We thus try to make a reconciliatory evolution model which could fit at least the present statistics of the multi-wavelength deep surveys. In this paper, we will first review the binary aggregation galaxy evolution model by Wang (1999) and by Wang \\& Biermann (2000) in section 2, where starburst/AGN activities may be triggered during the merger process, as well as an infrared burst phase from gas-rich mergers around a redshift of one. We will discuss the SED template we adopt in our calculation for the nearby starburst galaxies and a possible strong evolution of the spectral energy distribution of the dusty starburst merging systems at $z\\sim 1$. We thus could further investigate whether the infrared burst phase from gas-rich mergers around redshift $z\\sim 1$ is sufficient to account for the strong evolution also detected by ISO and submillimetre deep surveys. One set of cosmological parameters, namely $H_{0}=50\\,km/s/Mpc, \\,\\Omega=0.3$ and $\\Lambda=0.7$ is adopted in the calculation. ", + "conclusions": "} The exact broad-band spectra of faint IR sources is still not well defined. Considering deep surveys at various IR/submm wavelengths would help to simultaneously constrain the evolution properties and the typical spectral energy distribution of such sources. We show in this section the comparison of our model prediction with the ISOCAM $15\\,\\mu m$ survey data, IRAS $60\\,\\mu m$, ISOPHOT $90\\,\\mu m$ and $170\\,\\mu m$ (FIRBACK Survey), as well as the SCUBA $850\\,\\mu m$ data. We also calculated the redshift distribution of these sources within a certain flux range and the cosmic infrared background level. Fig.~\\ref{fig1} to Fig.~\\ref{fig3} are the model predictions for the European Large Area ISO Survey (ELAIS). This survey covered 12 $deg^2$ of the sky in four main areas and was carried out with the ISOPHOT instrument onboard the Infrared Space Observatory ISO, which is at least an order of magnitude deeper than the IRAS $100\\,\\mu m$ survey. It therefore provides an important constraint for our model of galaxy evolution. The majority of the optical identification of the detected sources are for interacting pairs or small groups of galaxies, which may indicate that the ELAIS sample includes a significant fraction of luminous infrared galaxies from galaxy mergers. Although there is some discrepancy in the data reduction, previous estimations show that the source counts are mostly in agreement with strongly evolving starburst models, with a rapid increase in the fraction of ULIGs towards high redshift (Efstathiou et al. 2000, Matsuhara et al. 2000, Serjeant et al. 2001). From our calculations, we see in Fig.~\\ref{fig1} to Fig.~\\ref{fig3} that the differential number counts of $90\\,\\mu m$, $15\\,\\mu m$ and $170\\,\\mu m$ for a reliable subset of the detected sources could be sufficiently accounted for by the infrared burst phase when a population of ultraluminous infrared sources with $L_{ir}>10^{12}\\,L_{\\odot}$ could be produced by the merger-triggered starburst/AGN activities at $z\\sim 1$ (Kawara et al. 1998, Elbaz et al. 1999, Efstathiou et al. 2000, Dole et al. 2001). The enormous infrared emission, especially at far-infrared wavelengths is modelled by a modified black body spectrum which we discussed in the previous section. We assume that the starburst merging system has a similar MIR emission feature as Arp220, but modified by the observed flux correlation of $S_{15}/S_{60}$ from IRAS and ISOCAM deep surveys. We found from the calculation that the dust temperature of these starburst merging systems would be higher than that of the nearby starburst ULIGs Arp220, with dust temperature $T= 65\\,K$ and $\\beta = 1.5$ for a best fit result. Fig.~\\ref{fig4} shows our model fitting for the differential counts of the IRAS $60\\,\\mu m$ deep survey, and Fig.~\\ref{fig5} shows the integrated number counts of submillimeter SCUBA deep survey at $850\\,\\mu m$ (Hacking et al. 1987, Moshir et al. 1992, Barger et al. 1999, Blain et al. 1999). Almost all the number counts could be reproduced quite well by such an evolutionary scenario, except for the ISOCAM $15\\,\\mu m$ differential number counts, where our model prediction shows a slight excess at the bright part of $S_{15}\\sim 2\\,mJy$. The reason could be that we simply adopt the mid-infrared emission feature of Arp 220 for the case of the starburst merging system around $z\\sim 1$. We hope to improve the current results by further theoretical modelling and additional observational constraints for the emission properties at mid infrared bands from future infrared missions. The infrared background in this calculation gives 2.4$nW\\,m^{-2}\\,sr^{-1}$ at 15$\\mu m$, 1.9$nW\\,m^{-2}\\,sr^{-1}$ at 60$\\mu m$, 3.8$nW\\,m^{-2}\\,sr^{-1}$ at 90$\\mu m$, 10.6$nW\\,m^{-2}\\,sr^{-1}$ at 170$\\mu m$, which are all consistent with current upper limits from TeV detections, COBE results and the resolved fraction of the CIRB by the deep ISO surveys (Funk et al., 1998, Guy et al. 2000, Hauser \\& Dwek 2001). The redshift distribution of the ISOCAM 15$\\mu m$ contributing sources within the detected flux range ($0.1\\,mJy\\sim 10\\,mJy$) from our model calculation is shown in Fig.~\\ref{fig6}. It gives a rough statistical finding that these luminous infrared sources cover a wide redshift range of $0.5\\sim 2.5$, peaking at $z\\sim 1$. Comparing our model prediction and the redshift distribution of 15$\\mu m$ sources with $S_{15}>120 \\mu Jy$ in the HDF North and the z-distribution of sources in the CFRS field (Flores et al. 1999, Cohen et al. 2000, Aussel et al. 2001), we found that the starburst mergers at $z\\sim 1$ in our model are good candidates for a strongly evolving population that results in the strong evolution in mid- and far-infrared deep surveys. Recent redshift estimation from sub-mm follow up of 10 known FIRBACK 170$\\mu m$ ISO sources by Scott et al. (2000) suggests that they are in a redshift range of $0\\sim 1.5$, still consistent with our current model predictions. However, these redshift determinations strongly depend on the assumption of the dust properties. We need further accurate measurements for the robust constraints of the models. We discussed in a previous paper that shifting the peak redshift of these ULIGs by a factor of 2 could affect the source count fitting of the IRAS 60$\\mu m$ deep survey, especially for a low redshift peak ($z<0.5$). Strong evolution of the ULIGs to $z\\sim 1$ may be the most reasonable case for the existing model constraints from both the infrared deep surveys and the cosmic infrared background upper limits from high energy TeV detections, as well as the indicated star formation history by UV/optical deep surveys (Lilly et al. 1996, Connolly et al. 1997, Madau et al. 1998). We plot the redshift distribution of ULIGs ($\\nu L_{\\nu}>10^{12}\\,L_{\\odot}$) to understand the evolutionary properties of the ULIGs from mergers in our model. A rapid increase in the number density of ULIGs up to $z\\sim 1$ is seen in Fig.~\\ref{fig7}, which is actually consistent with a scenario where galaxy merger rates increase dramatically during that epoch as seen from various observations and theoretical considerations (Zepf \\& Koo 1989, Carlberg 1992, Burkey et al. 1994, Carlberg et al. 1994). However, the number density of ULIGs decreases beyond $z\\sim 2-3$, which may reflect a stage when merger pairs are mostly dwarfs and the infrared emissions are less than $10^{12}\\,L_{\\odot}$ even with intensive starburst activities triggered by mergers. In this scenario, an infrared luminous tail of the luminosity function may form at $z\\sim 1$, with enormous infrared emission enhancement. There is still no firm statistical basis for the classification of starbursts and AGNs from current spectroscopies. We thus adopt the observed AGN local luminosity function of Rush et al. (1993) as a model constraint, and assume in our calculations that the observed starburst galaxies and Seyferts follow the same evolutionary track, based on naive thinking that starburst/AGN may both be triggered by galaxy interactions. We know the subtle differences in the dust emission properties could result in a different fraction of their contribution. This is still far too early to discuss here. We thus adopt only the SED of the Cloverleaf quasar which represents a phase poor in cold gas, as well as the dust enshrouded phase of F10214+4724 as two typical AGN templates in our calculation. We know from the result that the AGN contribution is only a small fraction of the whole and our current model prediction is within the present understanding of this issue, i.e. the starburst powered ULIGs are dominating over the AGN powered ones (Fig.~\\ref{fig7}) and may take over at higher redshifts and in the higher luminosity case (Lutz et al. 1998, Tran et al. 2001). \\begin{figure} \\label{fig1} \\includegraphics[height=8cm,width=8cm]{MS1717f1.eps} \\caption{The model prediction of the differential number counts of ISOCAM $15\\,\\mu m$ normalized to the Euclidean law ($dN/dS\\,*\\,S^{2.5}$). The data points are the normalized counts from a variety of ISO deep surveys (Elbaz et al. 1999). The line represents the sum of the contribution from three populations (starburst galaxies, spiral galaxies and Seyferts). In our model, we assume that the population of starburst galaxies, especially with spheroidal morphology, are the product of galaxy mergers which would experience infrared emission enhancement because of the merger-triggered starburst activities. Their contribution to the ISOCAM $15\\,\\mu m$ deep survey is shown by the dot-dashed line from our Monte-Carlo simulation. The dotted line corresponds to the non-evolving spiral galaxies, and the short dashed line is from Seyferts, which are assumed to have the same evolutionary track as starburst galaxies in a galaxy evolutionary scheme with galaxy mergers.\\label{fig1}} \\end{figure} \\begin{figure} \\label{fig2} \\includegraphics[height=8cm,width=8cm]{MS1717f2.eps} \\caption{The fitting of the ELAIS differential source count at $90\\,\\mu m$. The data are from C-90 filter of the C100 ISOPHOT detector array [filled squares: Efstathiou et al. 2000; open trianguler: Linden-V$\\phi$rnle et al. (2000); star: Juvela et al. (2000)]. The meaning of the lines is the same as in Fig.\\ref{fig1} \\label{fig2}} \\end{figure} \\begin{figure} \\label{fig3} \\includegraphics[height=8cm,width=8cm]{MS1717f3.eps} \\caption{The result of FIRBACK $170\\,\\mu m$ ISO deep survey differential number count fitting from our model calculation. The data are from Dole et al. (2001). \\label{fig3}} \\end{figure} \\begin{figure} \\label{fig4} \\includegraphics[height=8cm,width=8cm]{MS1717f4.eps} \\caption{The fitting of the IRAS $60\\,\\mu m$ source counts from three major infrared emitters (starburst galxies, spiral galaxies and Seyferts). The source counts of starburst galaxies and Seyferts are from the Monte-Carlo simulation where the evolution of both activities are triggered by galaxy-galaxy interactions/mergers during structure formation. The spiral galaxy is assumed to have a mild constant star formation history, i.e. a non-evolving population in our calculation. The data are from the IRAS Point Source Catalogue (1985)(PSC), Hacking et al. IRAS deep survey (HCH), FSC from deep surveys by Moshir et al. (1992) and Saunders (1990). \\label{fig4}} \\end{figure} \\begin{figure} \\label{fig5} \\includegraphics[height=8cm,width=8cm]{MS1717f5.eps} \\caption{Integral number counts at $850\\,\\mu m$. The open squares are from Blain et al. (1999), and the filled squares are from Barger, Cowie and Sanders (1999). The meaning of lines are the same as in previous figures. \\label{fig5}} \\end{figure} \\begin{figure} \\label{fig6} \\includegraphics[height=8cm,width=8cm]{MS1717f6.eps} \\caption{The redshift distribution of three infrared contributors (starburst galaxies, spiral galaxies and AGNs) at a flux range of $0.1\\,mJy\\sim\\, 10\\,mJy$ at $15\\,\\mu m$ from our model calculation. The redshift of these ISO far-infrared sources would cover a wide range and peak near $z\\sim 1$. We include also the AGN contribution in our model based on the observed AGN Local Luminosity Function from Rush et al. (1993). \\label{fig6}} \\end{figure} \\begin{figure} \\label{fig7} \\includegraphics[height=8cm,width=8cm]{MS1717f7.eps} \\caption{The redshift distribution of ultraluminous infrared sources (starbursts or AGNs) with $L_{ir}>10^{12}\\,L_{\\odot}$ from our model calculation. The number density of the ultraluminous infrared sources increase dramatically untill $z\\sim 1$, but quickly decrease afterward. \\label{fig7}} \\end{figure}" + }, + "0112/astro-ph0112256_arXiv.txt": { + "abstract": "Using the STIS spectrograph on \\emph{HST} we have obtained a grid of [\\ion{O}{3}]$\\lambda\\lambda$4959,5007 and H$\\beta$ emission-line spectra at $0\\farcs05\\times0\\farcs19$\\arcsec\\ and 60 km~s$^{-1}$ (FWHM) resolution that covers much of the NLR of NGC 1068. We find emitting knots that have blueshifted radial velocities up to 3200 km s\\( ^{-1} \\) relative to galaxy systemic, are \\( 70-150 \\) pc NE of the nucleus and up to 40 pc from the radio jet, emit several percent of the NLR line flux but no significant continuum, span a small fraction of the sky as seen from the nucleus, coincide with a region of enhanced IR coronal-line emission, show gradients in radial velocities of up to 2000 km~s\\( ^{-1} \\) in 7 pc, span velocity extents averaged over \\( 0\\farcs 1\\times 0\\farcs 2 \\) regions of up to \\( 1250 \\) km~s\\( ^{-1} \\), have ionization parameter \\( \\mathcal{U} \\ga 0.1 \\), and ionized masses \\( \\sim 200 \\) M\\( _{\\odot }/n_{e,4} \\) (\\( n_{e,4}=10^{4} \\) cm\\( ^{-3} \\)). The brightest parts of the blueshifted knots are often kinematically contiguous with more massive clouds nearer the jet that are moving with velocities of \\( \\leq 1300 \\) km s\\( ^{-1} \\) relative to galaxy systemic. However, some knots at \\( 1\\farcs 5-2\\farcs 5 \\) radii appear as bright points in a broken shell of radius \\( \\sim 0\\farcs 55 \\) (40 pc) that is expanding at up to 1500 km~s\\( ^{-1} \\), implying a dynamical age of \\( \\sim 1.3\\times 10^{4} \\) yrs. Between 2\\farcs5--4\\farcs5 from the nucleus, emission is redshifted relative to systemic, a pattern that we interpret as gas in the galaxy disk being pushed away from us by the NE radio lobe. We argue that the blueshifted knots are ablata from disintegrating molecular clouds that are being photoionized by the AGN, and are being accelerated radiatively by the AGN or mechanically by the radio jet. In their kinematic properties, the knots resemble the associated absorbers seen projected on the UV continua of some AGN. ", + "introduction": "Emission-line profiles that extend over several thousand km~s\\( ^{-1} \\) are hallmarks of activity in galaxy nuclei of all luminosities. Some AGN also show absorption lines at UV rest wavelengths, originating from warm gas located outside the region that generates the nonstellar continuum. This gas is invariably blueshifted over a range of several thousand km~s\\( ^{-1} \\) from the galaxy systemic velocity, yet often shows discrete components that each span several hundred km~s\\( ^{-1} \\) \\citep{ha97}. What process accelerates the clouds? Spatially complete spectral maps of nearby active galaxies can separate clouds that are agitated mechanically by the AGN from those that are simply illuminated by AGN photons (e.g.\\ undisturbed gas in an {}``ionization cone''). In this paper we report on observations made with the Space Telescope Imaging Spectrograph (STIS) and its medium-resolution (M) gratings to map the ``Narrow-Line Region\" (NLR) of the Seyfert galaxy NGC 1068. In this NLR, \\cite{al83} sampled a cloud complex that spans \\( \\sim 200 \\) pc, is centered \\( \\sim 210 \\) km~s\\( ^{-1} \\) blueward of galaxy systemic velocity, and has \\( \\sim 1670 \\) km~s\\( ^{-1} \\) FWHM (\\( \\sim 2200 \\) km~s\\( ^{-1} \\) according to the \\emph{HST} FOS spectra of \\cite{ca91}) velocity extent. The complex of clouds discovered by \\cite{wa68}, and mapped in detail by \\citet[CBT hereafter]{ce90} and \\cite{ar96} (see also \\citet{ba87}), was shown from FOS spectra \\citep{kr98} to be mostly photoionized by nuclear radiation, and to have bulk motion consistent with acceleration that is independent of radius \\citep[CKb hereafter]{CK00}. A recent STIS spectrum \\citep[KC1 hereafter]{kr00} shows that emission-line centroids near the {}``continuum hotspot'' (a feature with substantial scattered nuclear light located \\( \\sim 0\\farcs 17 \\) {[}12 pc{]} N of the optically obscured nucleus) exhibit a trend towards greater blueshifted velocities as the ionization potential of the line species increases, a result that extends the earlier finding by \\cite{mo96} who found that coronal, near-IR lines are extended across a \\( \\leq 4\\arcsec \\) region. Section 2 describes the acquisition and reduction of the STIS spectra. We show in \\S3 that the fastest knots have been accelerated to radial velocities blueshifted from galaxy systemic by $>$3200 km~s\\( ^{-1} \\), a pattern not evident in the single STIS spectrum discussed by KC1 and CKb. We consider various mechanisms capable of accelerating clouds and knots in \\S4, outline why the knots may be related to quasar ``associated absorbers\" in \\S5, briefly consider future observations to test this hypothesis in \\S6, and summarize our conclusions in \\S7. In this paper, we assume a distance to NGC 1068 of 14.4 Mpc, so 1\\arcsec = 70 pc. ", + "conclusions": "While the clouds might have been as blended at \\emph{HST} resolution as they are from the ground, in \\S3 we showed that our deep spectra have resolved spatially the NLR. We have also resolved much of it kinematically, although not down to the thermal width (17 km~s\\( ^{-1} \\) in \\( 10^{4} \\) K gas) in part because of inadequate spectral resolution and in part because several components are still projected along each line of sight and may have turbulent substructure. To track the energy flow and to isolate possible shock-excited structures, we must now constrain the internal properties and space velocities of the NLR clouds. \\subsection{\\label{sec:photo}Photoionization of the NLR Clouds} The NLR clouds are bright features embedded in the large-scale ionization cones and are distributed over several arc-seconds, so have been studied extensively from space and, in aggregate, from the ground. FOS spectroscopy lacked spatial resolution to separate clouds from their surroundings. Subsequently, KC obtained a STIS spectrum along one cut through the NLR using low-resolution gratings that spans \\( \\lambda \\)\\( \\lambda \\)115 -- 1027 nm. Along the slit, they analyzed continuum fluxes (CKa), line fluxes (KC1), and gas kinematics (CKb). In KC1 they reproduced the observed flux ratios with AGN photoionization models that require considerable absorption along sight-lines from the nucleus that intersect the galaxy gas disk. They confirmed a correlation between ionization potential and line blueshift from galaxy systemic velocity that \\cite{mo96} had discovered in their ground-based IR spectra. KC1 sampled several clouds NE of the nucleus, so their photoionization models also address the new features that we have mapped. Their models combine matter- and ionization-bounded clouds, the former arising plausibly from low-density, photoevaporated envelopes of the latter \\citep{bi96}. Matter-bounded clouds may filter AGN photons that impinge on the ionization-bounded cores, so internal dust would play a critical role in setting their emission-line flux ratios \\citep{DoGr01,gr01}. The models of KC1 required that the blue-wing clouds be optically thin above the Lyman limit, and need another photon source between 1\\farcs4-1\\farcs8 to explain strong NUV lines and a larger \\ion{He}{2}\\( \\lambda \\)4686/H\\( \\beta \\) ratio along 202\\arcdeg\\ P.A. CKb suggested that clouds are decelerating here, forming a \\( \\sim \\)1000 km~s\\( ^{-1} \\) shock that would brighten high-ionization lines as observed. They lacked the sensitivity and spatial coverage of \\cite{mo96} to coronal-line emission, so unlike those authors found no such emission beyond cloud B and the UV continuum hotspot. In CKa they established that the scatterer is tenuous and coincides with the deceleration region. They found reasonable fits to blue-wing line fluxes if 85--90\\% of the flux comes from tenuous gas with \\( (\\log \\mathcal{U}, \\)\\( n_{H},\\log N_{H})=(-1.45\\pm 0.15,3\\times 10^{3}\\, \\rm{cm^{-3}},21.18\\pm 0.18\\, \\rm{cm^{-2}}) \\) and the rest from denser clouds with \\( (\\log \\mathcal{U} \\)\\( ,n_{H},\\log N_{H})=(-2.9,2\\times 10^{4}\\, \\rm{cm^{-3}},20.7\\pm 0.1\\, \\rm{cm^{-2}}) \\), both components needing dust fractions of 10--25\\% (larger values closer to nucleus) to suppress Ly\\( \\alpha \\) and C IV to their observed dereddened fluxes. KC1 found that the red-wing clouds have the higher dust fraction (25\\%) and are exposed to a lower ionizing flux (especially above the \\ion{He}{2} Lyman limit), indicating absorption of AGN photons on the way to the red wing gas, and thence by the tenuous blue-wing component along our line of sight. \\cite{mo96} find IR coronal lines across the inner 4\\arcsec\\ diameter that emit out to 2000 km~s\\( ^{-1} \\) blueshift from systemic velocity, and which can arise either in gas photoionized by a hard UV continuum or in a hot (\\( \\sim 10^{6} \\) K) collisionally ionized (perhaps shocked) plasma. \\citet{ol01} note that the {[}\\ion{Fe}{2}{]} emission is unlikely to come from shocks because they find that the flux ratio {[}\\ion{Fe}{2}{]}/{[}\\ion{P}{2}{]}$\\sim1$, 5\\% of the value seen in supernova remnants. X-ray results also support photoionization NE of the nucleus, beyond radius $r_i$. First, \\cite{Yo01} show that filaments in the \\emph{Chandra} ACIS image correlate well with those in [\\ion{O}{3}]$\\lambda$5007 and H$\\alpha$+[\\ion{N}{2}] images. Second, the \\emph{Newton/XMM} RGS spectrum reported by \\cite{Pa00} \\cite[also][]{be01} show that gas in this region is photoionized, based on the detection of narrow radiative recombination continua, line flux ratios within the helium-like triplets, and the weakness of Fe L-shell compared to K-shell emissions. \\subsection{\\label{sec:velocities}Deprojected Velocity Field} Ground-based Fabry-Perot (CBT) and integral-field spectra (PFBW) have constrained the bulk motion and space distribution of the NLR clouds. CBT found an overall distribution of clouds in a thick bicone with maximum half opening angle \\( \\sim 41\\arcdeg \\) and whose axis is inclined \\( \\sim 85\\arcdeg \\) to our line of sight; PFBW showed that gas coincident with the NE radio lobe beyond 2\\farcs5 radius is predominantly redshifted. CKb also posited a biconical flow to match the velocity centroids along their single STIS slit, requiring outflow velocities that reach 1300 km~s\\( ^{-1} \\) at $r_i\\sim1\\farcs7$ NE of the nucleus and then decrease beyond. Limited spatial coverage of a structure whose symmetry axis lies near the plane of the sky meant that their kinematical data are consistent with either an outflow that is radial from the nucleus or one that is roughly perpendicular to the jet axis. The archival M-grating spectrum (Fig.\\ \\ref{fig:velocities}) shows that knots HV6 and 7 are comparable to the largest blueshifts in the model geometry of CKb. \\placefigure{\\ref{fig:velocities}} We see abrupt jumps in velocity as the STIS slits cross the clouds and knots evident in the FOC narrow-band image. In Fig.\\ \\ref{fig:fig5}, we compare our line profiles with the projections of several models of radial and cylindrical outflow. By increasing the half angle of the NE bicone from \\( \\sim \\)40 to 50\\arcdeg\\ within the radius \\( r_{i} \\) of the CKb model where outflow velocities begin to decrease, we better match the trend in the centroids of radial velocities of the blueshifted features in our spectra. The velocity limits of our two models are shown in the Figure, and correspond to: (solid line) increasing the maximum outflow to at least 2500 km~s\\( ^{-1} \\) at $r_i$, and (long dashed line) having the gas expand from the jet axis with all of its motion along our line of sight. We also plot the outflow model of CKb in short dashed lines. The gas becomes increasingly blueshifted (on average) out to radius 2\\arcsec, implying its acceleration on that scale. However, the observed gas does not fill all of the NE bicone, because redshifted emission is observed at much lower velocities than predicted by all of the outflow models until the NE radio lobe is reached. Emission then jumps to a range that extends from zero to $>1000$ km~s$^{-1}$ redshift (relative to systemic velocity) and persists from 3\\arcsec\\ to the middle of the NE radio lobe (which occurs at the top of the panels in Fig.\\ 2); the emitting features are spatially compact. KC1 noted that the emission-line flux ratios of the red wing gas are consistent with irradiation by a heavily absorbed component. If these clouds are accelerated radiatively, they may not attain velocities as high as those that were irradiated directly. \\placefigure{\\ref{fig:fig5}} \\placefigure{\\ref{fig:fig6}} As suggested by \\cite{te01} based on infrared [\\ion{Fe}{2}] spectra and PFBW based on spatially deconvolved optical spectra, the red wing gas beyond 2\\farcs5 radius is apparently being pushed into the dense galactic disk by the lateral expansion of the NE radio lobe. Our favored scenario is shown in Fig.\\ \\ref{fig:fig6}. The jet is inclined to the disk. Eventually, at radius 2\\arcsec, it bursts out of the disk and inflates the radio lobe. The lobe expands laterally in all directions, but is only visible where it pushes into the disk and creates the zero to redshifted emission. The large blueshifted extensions on the high-velocity clouds suggest that the clouds are moving slower than the HV knots. The HV flow must have been accelerated relative to the clouds or reoriented closer to our line of sight by a combination of bouyancy, radiative acceleration, and lobe/cloud collisions \\citep{ta92}. Depending upon their dynamics, the larger clouds may lag the knots because of larger drag forces, as KC1 have proposed for the red wing component. \\subsection{Acceleration of Clouds \\& Knots} \\subsubsection{\\label{sec:jet}Jet Interactions with Clouds C \\& D} Using the long-slit mode of the FOC, \\cite{Ax97} found that line profiles are double-peaked and have highest gaseous excitation across the radio jet near clouds G and what we call K. On this basis they argued that hot shocked gas in the jet cocoon expands clouds from the jet axis. Nearby, \\cite{kr98} found evidence from emission-line ratios in an FOS spectrum for cosmic-ray heating of the emission-line gas. The Merlin and VLA images of \\citet[G96 hereafter]{Ga96}, when registered using the astrometry of \\cite{ca97} (see red contours in Fig.\\ 1), place NLR clouds B, C, D, and F adjacent to the radio jet; in particular, clouds B and C straddle a stationary radio knot \\citep{ro01} where the jet bends \\( \\sim 23 \\)\\arcdeg. \\cite{bi98} have modeled this complex as a jet-cloud interaction. What constraints do our data impose? KC1 found that cloud C is photoionized. However, decomposing its profile in Fig. \\ref{fig:f12} into Gaussians isolates a component centered \\( \\sim 300 \\) km s\\( ^{-1} \\) blueward of systemic and a wing component centered 300 km s\\( ^{-1} \\) to the red, both with dispersions \\( \\sim 500 \\) km s\\( ^{-1} \\). The blue component has flux ratio {[}\\ion{O}{3}{]$\\lambda5007$}/H\\( \\beta \\sim 16 \\) \\emph{vs.} \\( \\sim 8-10 \\) for the red component. While the considerable line widths are consistent with shocks, the flux ratios are consistent with photoionization by the AGN with the red wing gas more absorbed. Thus, deflection of the jet at cloud C is not clearly connected to the excitation of the cloud. Cloud E, which is farther from the jet than cloud C, shows a similar blue feature but not one in its red wing. The NE component of the nuclear radio triple is 0\\farcs75 NE of knot S1 (the nucleus). G96 note that this jet component, like the radio knot near cloud D, has a flat synchrotron spectrum and extended emission that is misaligned with the jet. However, the jet does not bend at the NE component nor is there maser emission, arguing against a shock front. It is therefore interesting that the emission-line profiles of cloud D are double peaked. The blue peak has flux ratio {[}\\ion{O}{3}{]}/H\\( \\beta \\sim 15 \\), but now the red peak ratio decreases monotonically through the profile from 17 to $<3$ over a 1000 km s\\( ^{-1} \\) range. This decline implies that gas density increases with the velocity deviation from galaxy systemic, not what would be expected if most of the cloud mass is at velocities close to rest in the galaxy. G96 note that a cloud interacting with the jet must be sufficiently massive to avoid acceleration, implying \\begin{eqnarray*} M_{cloud\\, B}\\geq 40\\, M_{\\bigodot }\\left( \\frac{r_{j}}{7.5\\, pc}\\right) ^{2}\\left( \\frac{\\rho _{j}}{5\\times 10^{-27}\\, g\\, cm^{-3}}\\right) \\times & & \\end{eqnarray*} \\begin{eqnarray*} \\left( \\frac{v_{j}}{15,000\\, km\\, s^{-1}}\\right) \\left( \\frac{t_{lobe}}{3.5\\times 10^{5}\\, yr}\\right). & & \\end{eqnarray*} Here jet density and velocity are normalized to the values appropriate for the lighter jet model of \\cite{Wi87} for the NE radio lobe, $t_{lobe}$ is the age of the NE radio lobe from spectral aging as estimated by G96, and the jet radius \\( r_{j} \\) was measured by \\cite{ga94}. If the average NLR cloud density is \\( \\sim \\)150 cm\\( ^{-3} \\), the entries in Table 2 show that clouds are massive enough to remain adjacent to the deflecting jet. This value is consistent with the mean densities of Galactic molecular clouds \\citep{so87}. Elsewhere, radio and line emission are uncorrelated. G96 note that this is expected if the jet has swept all gas from the region. The high-velocity knots might then be ablated from molecular clouds in the galaxy disk that have rotated into the jet. The clouds accelerated would all be blueshifted because they would move predominantly away from and above the galaxy disk. Similar clumps below the jet would sink into the denser disk where they would be subject to stronger drag forces and hence decelerated more effectively (consistent with the gas distribution in the spectral maps of PFBW). \\subsubsection{Shocks Near the X-ray Subpeak?}\\label{vshape} Three characteristics of the \\emph{XMM/Newton} X-ray spectrum of NGC 1068 within 20\\arcsec\\ of the nucleus indicate that the gas is predominantly photoionized: (i) narrow radiative recombination continua; (ii) in the He-like triplets the forbidden lines are stronger than the resonance lines; and (iii) Fe L emission is weak compared to the K shell emission of the abundant elements \\citep{ki01,be01}. However, subtraction of the predictions of a photoionization model from the spectrum shows residuals that may indicate a collisionally excited component or other effects \\citep{Pa00}. Moreover, near the V-shaped line-emitting feature whose apex is 3\\arcsec\\ NE of the nucleus (at top of Fig.\\ 1), the \\( 4-26 \\) \\AA\\ spectrum from the \\emph{Chandra} HETGS \\citep{og01} shows that resonance lines are brighter than forbidden lines, implying either a collisionally ionized component or strong resonant scattering. \\cite{og01} show that the collisional component lies at the X-ray subpeak reported by \\cite{Yo01}. The X-ray subpeak may represent gas in the galactic disk on the far side of the NE lobe and being compressed by it (cf.\\ PFBW). As discussed in \\S\\ref{sec:photo}, photoionization models of the UV and optical STIS spectra of KC1 also support this additional power source near the subpeak. \\subsubsection{\\label{sec:shocks}Shocks Near Clouds G-K?} Two lines of evidence argue for shocks near clouds G-K: \\begin{enumerate} \\item Clouds G and K are close to the ``NE radio subpeak\", located 2\\farcs1 NE of the obscured nucleus \\citep{Wi87}, suggesting dissipation of kinetic energy and possibly cloud acceleration by the radio jet. \\item Nearby we find the Doppler ellipsoid described in \\S\\ref{sec:cloudinfo}, whose expansion velocity of 1500 km~s\\( ^{-1} \\) will generate a post-shock temperature of \\( \\sim 3\\times 10^{7} \\) K if the ambient medium is relatively stationary, fully ionized, and has 10\\% He abundance. \\end{enumerate} To estimate the mechanical luminosity of the Doppler ellipsoid, we assume that energy is injected at constant rate, that the shell shock is strong so that we can ignore the thermal pressure of the ambient gas, and that the bubble is within the scale height of the galaxy disk so that the ambient ISM density is constant. Then in the energy conserving phase of the bubble interior, the bubble radius and expansion velocity are related: \\( R=1.67vt. \\) By observation, \\( t=0.012 \\) Myr. Mechanical luminosity \\( L \\) follows from \\( R(t)=31(L_{36}/n_{0})^{0.2}t_{6}^{0.6} \\) pc = 100 pc with \\( t_{6}=0.012 \\) (where \\( L_{36} \\) has units \\( 10^{36} \\) ergs s\\( ^{-1} \\) and \\( n_{0} \\) is the density of the ambient ISM in cm\\( ^{-3} \\)), so \\( L=5.2\\times 10^{7}n_{0} \\) L\\( _{\\odot } \\). Overall, the filaments that delineate the bubble and maybe its associated ISM shock emit $<$2\\% of the NLR {[}\\ion{O}{3}{]} and H\\( \\beta \\) fluxes. \\subsubsection{\\label{sec:accel}Radiative Acceleration of Dusty Clouds} Dust absorption dominates opacity in a photoionized plasma like this NLR when ionization parameter \\( \\mathcal{U} >\\alpha (T_{e})/c\\kappa \\), where \\( c \\) is lightspeed, \\( \\kappa \\sim 10^{-21} \\) cm$^2$ atom$^{-1}$ is the opacity assuming the standard Galactic reddening curve, and \\( \\alpha (T_{e})\\sim 2\\times 10^{-13} \\) is the recombination coefficient. The dust will compete successfully for ionizing photons in the highly ionized zone, suppressing line flux. A strong pressure gradient will also be set up so that the plasma density will be much enhanced by the time lower ionization species emit. In consequence, a lower \\( \\mathcal{U} \\sim 0.007 \\) will provide an apparent model fit to the spectrum, close to the value preferred by KC for their denser component. This ionization parameter is closely related to the value of \\( \\mathcal{U} \\) at which radiation pressure starts to dominate either the pressure gradient, or the dynamical acceleration of the ionized plasma. Thus, the flow of dusty clouds falling toward the nucleus is intrinsically unstable above the dust sublimation point, with clouds being driven into outflow in the direction of slowest accretion (i.e.\\ within the ionization cones). \\cite{DoGr01} detail this scenario. In hydrodynamical simulations of cloud acceleration \\citep{sc95}, tenuous gas streams ablate from surfaces of massive clouds, which may be the source of the high-velocity features we observe. In this NLR, the flux in the ionizing continuum needed to maintain the observed line luminosity (\\( \\sim 3\\times 10^{44} \\) ergs~s\\( ^{-1} \\), KC1) is comparable to the FIR luminosity 1.5\\( \\times 10^{11} \\) L\\( _{\\odot } \\). Dusty clouds are accelerated radiatively by \\( \\sim 10^{-5} \\) cm~s\\( ^{-2} \\), so can reach their observed mean velocities of \\( \\sim 500 \\) km~s\\( ^{-1} \\) over 35 pc. In Fig.~3, the cloud line profiles indeed show less velocity structure than do the knots. This scenario ignores the role of the radio jet in agitating the NLR clouds \\citep{bi98}. But, as discussed earlier, its influence may be localized to line broadening and ablation in specific jet/cloud and lobe/cloud interactions, for example between B and C \\citep{Ga96} and G and H \\citep{Ax97}. While the region of collisional line emission coincides with the base of the NE radio lobe, numerical simulations of a breakout lobe do not show (D.\\ Balsara, private communication) a strong backflow vortex at its base that might form high-velocity shocks. In any event, many of the high-velocity knots have similar line profiles, and are found \\( >40 \\) pc from the brighter parts of the jet. Our STIS spectral maps have resolved the spatio-kinematic structure of the NLR in NGC 1068, and reveal that many compact knots in the FOC image have large, exclusively blue-shifted velocities, ranging up to $\\sim3200$ km~s$^{-1}$ from systemic. If these knots are optically thick in the UV continuum, they are good candidates for ``associated absorber\" clouds present in other AGN. Lying 70--150 pc from the nucleus, several form a broken spatio-kinematic ring of diameter 35 pc that is expanding at up to 1500 km~s\\( ^{-1} \\). The ring is adjacent to the more massive but slower moving NLR clouds G--K, so may originate from ablata streams from disintegrating molecular clouds that are being photoionized and accelerated radiatively by the AGN or mechanically by a knot in the radio jet. By resolving this NLR, we have demonstrated that ground-based integral-field spectrometers will be able to constrain the properties of the emitting gas in NGC 1068 by using classical plasma diagnostics in the optical and near-IR. More fundamentally, we may have a convenient view onto gas associated with an important component of AGN that is hard to study elsewhere because it is usually seen in absorption." + }, + "0112/astro-ph0112126_arXiv.txt": { + "abstract": "Near-infrared (NIR) spectra of the subluminous Type~Ia supernova SN~1999by are presented which cover the time evolution from about 4 days before to 2 weeks after maximum light. Analysis of these data was accomplished through the construction of an extended set of delayed detonation (DD) models covering the entire range of normal to subluminous SNe~Ia. The explosion, light curves (LC), and the time evolution of the synthetic spectra were calculated self-consistently for each model with the only free parameters being the initial structure of the white dwarf (WD) and the description of the nuclear burning front during the explosion. From these, one model was selected for SN~1999by by matching the synthetic and observed optical light curves, principly the rapid brightness decline. DD models require a minimum amount of burning during the deflagration phase which implies a lower limit for the $^{56}Ni$ mass of about $0.1 M_\\odot$ and consequently a lower limit for the SN brightness. The models which best match the optical light curve of SN~1999by were those with a $^{56}Ni$ production close to this theoretical minimum. The data are consistent with little or no interstellar reddening ($E(B-V) \\leq 0.12^m$) and we derive a distance or $11 \\pm 2.5$ Mpc for SN~1999by in agreement with other estimates. Without any modification, the synthetic spectra from this subluminous model match reasonably well the observed IR spectra taken on May 6, May 10, May 16 and May 24, 1999. These dates correspond roughly to $-4$~d, 0~d, and 6~d and 14~d after maximum light. Prior to maximum, the NIR spectra of SN~1999by are dominated by products of explosive carbon burning (O, Mg), and Si. Spectra taken after maximum light are dominated by products of incomplete Si burning. Unlike the behavior of normal Type~Ia SNe, lines from iron-group elements only begin to show up in our last spectrum taken about two weeks after maximum light. The implied distribution of elements in velocity space agrees well with the DD model predictions for a subluminous SN~Ia. Regardless of the explosion model, the long duration of the phases dominated by layers of explosive carbon and oxygen burning argues that SN~1999by was the explosion of a white dwarf at or near the Chandrasekhar mass. The good agreement between the observations and the models without fine-tuning a large number of free parameters suggests that DD models are a good description of at least subluminous Type~Ia SNe. Pure deflagration scenarios or mergers are unlikely and helium-triggered explosions can be ruled out. However, problems for DD models still remain, as the data seem to be at odds with recent 3-D models of the deflagration phase which predict significant mixing of the inner layers of the white dwarf prior to detonation. Possible solutions include the effects of rapid rotation on the propagation of nuclear flames during the explosive phase of burning, or extensive burning of carbon just prior to the runaway. ", + "introduction": "On April 30, SN~1999by was independently discovered at about $15^m$ by \\citet{arbor99}, and \\citet{papenkova99}. Images with the Katzman Automatic Imaging Telescope \\citep{treffers97,li99} of the same field provide an upper limit of $19.3^m$ on April 25.2 UT. The supernova (SN) was found in the Sb galaxy NGC~2841 which has been the host of three previous supernovae: SN~1912A, SN~1957A, and SN~1972R \\citep{papenkova99}. Based on optical spectra, SN~1999by was identified as a Type~Ia SN \\citep{gerardy99}. Shortly thereafter, \\citet{garnavich99} reported that SN~1999by showed stronger than normal \\ion{Si}{2} 5800~\\AA\\ absorption and depressed flux near 4000~\\AA, suggesting that SN~1999by would be a significantly subluminous Type~Ia event. According to \\citet{bonanos99}, SN~1999by reached a maximum light of $m_B=13.80^m \\pm 0.02 ^m$ on UT 1999 May 10.5, and a maximum in the $V$ band of $m_V=13.36^m \\pm 0.02 ^m$. The Lyon/Meudon Extragalactic Database (LEDA)\\footnote{http://www-obs.univ-lyon1.fr/leda/home\\_leda.html}, gives the heliocentric radial velocity of NGC~2841, corrected for Virgo infall, as $811.545$ km~s$^{-1}$. Using $H_0=65$ km~s$^{-1}$~Mpc$^{-1}$ puts the distance to NGC~2841 at 12.5 Mpc with a distance modulus of $30.5^m$. Using this distance modulus and the photometry of \\citet{bonanos99}, the absolute peak magnitudes of SN~1999by are $M_B=-16.68^m$ and $M_V=-17.12^m$. Thus, SN~1999by is underluminous by roughly 2.5 magnitudes as compared to a typical Type~Ia supernova (SN~Ia). Detailed analysis of the optical light curves (LCs) by \\citet{toth00} confirm these basic results. The light curve of SN~1999by shows a very steep post-maximum brightness decline $M_V(\\Delta M_{\\Delta t=15d})$ of $1.35^m$ to $1.45^m$. Based on detailed fits of the LCs, \\citet{toth00} find the interstellar reddening to be $E(B-V) \\leq 0.1^m$ which is in agreement with the values for galactic reddening given by \\citet[$E(B-V)=0.015^m$]{schlegel98}, and \\citet[$E(B-V) \\approx 0^m$]{burstein82}. Recently, \\citet{garnavich} provided detailed optical LCs and spectra and found values consistent with previous data. These measurements imply that SN~1999by was as underluminous as SN~1991bg, the prototypical example of the subluminous Type~Ia subclass \\citep{filippenko92,leibundgut93}. Other members include SN~1992K \\citep{hamuy96a,hamuy96b}, SN~1997cn \\citep{turatto98}, and SN~1998de \\citep{modjaz00}. Some defining characteristics of the subclass are rapidly declining light curves ($M_B(\\Delta M_{\\Delta t=15d}) \\simeq 1.9^m$), peak magnitudes fainter than normal by 2--3 mag, and redder colors at maximum light ($(B-V)_{\\rm max}$ $\\simeq$ 0.4--0.5$^m$). One of the currently active areas in SNe~Ia research concerns the nature of these subluminous events. Theoretical interpretations of subluminous SN~Ia include all three types of explosion mechanism: the centrally triggered detonation of a sub-Chandrasekhar mass WD, deflagration and delayed detonations of massive WD and two merging WDs (see sect. 3 \\& Woosley \\& Weaver, 1994 ; H\\\"oflich et al., 1995; Nugent et al., 1995, Milne et al. 1999). The possibility has also been raised that SN1991bg-like SNe~Ia should not be classified with other SNe~Ia at all as they may arise from different progenitors. SN~1999by is one of the best observed SNe~Ia with data superior to that of SN~1991bg. In addition to the studies of optical light curves and spectra mentioned above, detailed polarization spectra of SN~1999by have been obtained and analyzed \\citep{howell01}. Whereas `normal' SNe~Ia tend to show little or no polarization \\citep[e.g.][]{wang97}, this supernova was significantly polarized, up to 0.7\\%, indicating an overall asphericity of $\\approx 20\\%$. This result suggests that there may be a connection between the observed asphericity and the subluminosity in SNe~Ia. In recent years it has become apparent that infrared spectroscopic observations can be used as valuable tools to determine the chemical structure of SN~Ia envelopes. For instance, near-infrared spectra can be used to locate the boundaries between explosive carbon and oxygen burning, or between complete and incomplete silicon burning by measuring \\ion{Mg}{2}, \\ion{Ca}{2} and iron-group lines \\citep{wheeler98}. However, the difficulties involved in obtaining high quality infrared spectra of supernovae has limited IR studies to a small handful of objects. The study of SN~1994D by \\citet{meikle96} is the only one to include both pre- and post-maximum spectra. Since most of their spectra was of the 1--1.3~\\micron\\ region, \\citet{meikle96} concentrated on determining the origin of a feature at 1.05~\\micron. They concluded that this feature might be due to either \\ion{He}{1} 1.083~\\micron\\ or \\ion{Mg}{2}~1.0926 \\micron, but found difficulties with both identifications. Detailed modeling of SN~1994D and SN~1986G was able to reproduce the basic NIR spectral features and their evolution \\citep{hoflich97,wheeler98}. In particular, the identification of the 1.05 \\micron\\ feature in SN~1994D as due to \\ion{Mg}{2} could be established (recently confirmed by \\citealt{lentz01}), and a broad feature, appearing between $\\approx 1.5$ and 1.9 \\micron\\ was identified as a blend of iron-group elements. We note that the nature of subluminous SNe~Ia is important for the use of SNe~Ia as distance indicators. In particular, the calibration of the brightness decline relation depends critically on subluminous SNe~Ia because they significantly increase the required baseline observations. From this perspective, it may turn out to be critical to determine whether all SNe~Ia form a homogeneous class of objects or not. However, current limitations of model calculations will not allow us to improve the accuracy of the current estimates for the absolute distance to SN~1999by, and the possible implications of our results on the use of SNe~Ia for cosmology is not the subject of this paper. In this paper, we present near-infrared spectra of SN~1999by covering the time evolution of the supernova from about four days before to two weeks after maximum light. Detailed models of the SN explosion based on a delayed detonation (DD) scenario are used to analyze the data. The explosion models, light curves, and synthetic spectra are calculated in a self-consistent manner. Given the initial structure of the progenitor and description of the nuclear burning front, the light curves and spectra are calculated from the explosion model without any further tuning. The observations and data reduction are discussed in \\S 2. In \\S 3, general properties of the explosion models, light curves and spectra are discussed for normal and subluminous SNe~Ia within the DD scenario. In \\S 4, a specific model for this SN is chosen by matching the properties of the predicted light curve to those observed in optical studies of SN~1999by. Synthetic optical and IR spectra for this model are then presented and the latter are compared to the observed IR spectra. Detailed line identifications for the infrared features are provided. The effects on the predicted spectra of large scale mixing during the deflagration phase also examined. In \\S 5, we discuss the implications for progenitors and alternative explosion mechanisms. We close in \\S 6 with a final discussion of the results for SN 1999by, and put our findings in context with other SNe~Ia and different explosion scenarios. \\section {Observations} Low-dispersion (R $\\approx$ 700), near-infrared spectra of SN~1999by from 1.0--2.4~\\micron\\ were obtained with the 2.4m Hiltner Telescope at MDM Observatory during the nights of 6--10 May 1999. The data were collected using TIFKAM (a.k.a.\\ ONIS), a high-throughput infrared imager and spectrograph with an ALLADIN 512 $\\times$ 1024 InSb detector. This instrument can be operated with standard \\textit{J}, \\textit{H}, and \\textit{K} filters for broadband imaging, or with a variety of grisms, blocking filters, and an east-west oriented $0\\farcs 6$ slit for low ($R \\approx 700$) and moderate ($R \\approx 1400$) resolution spectroscopy. SN~1999by was observed with three different spectroscopic setups, which covered the 0.96--1.80~\\micron, 1.2--2.2~\\micron, and 1.95--2.4~\\micron\\ wavelength regions. The observations were broken into multiple 600~s exposures. Between each exposure, the supernova light was dithered along the slit to minimize the effect of detector defects and provide first-order background subtraction. Total exposure times vary from spectrum to spectrum and are listed in the log of NIR observations in Table~\\ref{nirlog}. Wavelength calibration of the spectra was achieved by observing neon, argon, and xenon arc lamps. The spectra were corrected for telluric absorption by observing A~stars and early G~dwarfs at similar airmass, chosen from the Bright Star Catalog \\citep{BSC}. Applying the procedure described by \\citet{hanson98}, stellar features were removed from the G~dwarf spectra by dividing by a normalized solar spectrum \\citep{solar1, solar2}\\footnote{NSO/Kitt Peak FTS data used here were produced by NSF/NOAO.}. The results were used to correct for telluric absorption in the A~stars. Hydrogen features in the corrected A~star spectra were removed from the raw A~star spectra and the results were used to correct the target data for telluric absorption. [For further discussion of this procedure see \\citealt{hanson98,HCR96}, and references therein.] The instrumental response was calibrated by matching the continuum of the A~star telluric standards to the stellar atmosphere models of \\citet{kurucz93}. The first three nights (6-8 May) were photometric, and \\textit{J}, \\textit{H}, and \\textit{K} broadband images of SN~1999by and \\citet{persson} photometric standards. The resulting photometry is listed in Table~\\ref{nirlog} and was used to set the flux levels of the corresponding spectra. A square bandpass was assumed and the flux levels thus attained are believed accurate to $\\sim$ 20--30\\%. For May 9 and 10 no flux information was available, but the observed spectral bands had large overlaps, and the relative flux levels were set by matching the data in the overlapping regions. Near-infrared spectra of a star $29\\farcs 5$ north and $11\\farcs 7$ west of SN~1999by \\citep{papenkova99} were also obtained and reduced in the same manner. This star was then used as a local relative spectrophotometric standard to reduce spectroscopic data from later observing runs. Spectra of SN~1999by covering 1.0--1.8~\\micron\\ were collected on five more nights during the following two weeks (16 May, 18 May, 20 May, 21 May and 24 May). The first three were obtained by P. Martini and A. Steed using TIFKAM on the 2.4m telescope at MDM, and the last two by S. Sakai using TIFKAM on the 2.1m telescope at the Kitt Peak National Observatory. Figure~\\ref{nirspec} shows a plot of all the NIR spectra. Since absolute fluxes were only available for the first three spectra, the data are presented in arbitrary flux units, and they have been shifted vertically for clarity. Regions of very low S/N due to strong telluric absorption have been omitted. The epochs listed are relative to 10 May 1999, the date of $V_{Max}$ given by \\citet{bonanos99}. Like all the subsequent plots, the data have been corrected to the 638 km~s$^{-1}$ redshift of NGC~2841 \\citep{devaucouleurs91}. \\section {Models for the Explosion, Light Curves, and Spectra} There is general agreement that SNe~Ia result from some process involving the combustion of a degenerate white dwarf \\citep{hoyle60}. Within this general picture, three classes of models have been considered: (1) An explosion of a CO-WD, with mass close to the Chandrasekhar limit, which accretes mass through Roche-lobe overflow from an evolved companion star \\citep{whelan73}. The explosion is triggered by compressional heating near the WD center. (2) An explosion of a rotating configuration formed from the merging of two low-mass WDs, caused by the loss of angular momentum through gravitational radiation \\citep{webbink84,iben84,paczynski85}. (3) Explosion of a low mass CO-WD triggered by the detonation of a helium layer \\citep{nomoto80,woosley80,woosley86}. Only the first two models appear to be viable for normal Type~Ia SNe as the third, the sub-Chandrasekhar WD model, has been ruled out on the basis of predicted light curves and spectra \\citep{hoflich96b,nugent97}. Still, theoretical interpretations of subluminous SN~Ia include all three types of explosion mechanism \\citep{woosley94,hkw95,nugent95,milne99}. The possibility has also been raised that the subluminous SN~1991bg-like supernovae may arise from different progenitors and should not be classified with other SNe~Ia at all. Within $M_{Ch}$ models, it is believed that the burning front starts as a subsonic deflagration. However, the time evolution of the burning front is still an open question. That is, whether the deflagration front burns through the entire WD \\citep{nomoto84}, or alternatively, transitions into a supersonic detonation mode as suggested in the delayed detonation (DD) model \\citep{khokhlov91,woosley94,yamaoka92}. DD models have been found to reproduce the optical and infrared light curves and spectra of `typical' SNe~Ia reasonably well (\\citealt{hoflich95}, H95 hereafter; \\citealt{hk96,hoflich96b,fisher98,nugent97,wheeler98,lentz01}). In addition, DD models provide a natural explanation for the brightness decline relation $M (\\Delta M_{\\Delta t=15d})$ observed in the light curves of SNe~Ia \\citep{phillips87,hamuy95,hamuy96b,suntzeff99}. This is a consequence of the temperature dependence of the opacity and the lack of a dependence of the explosion energy on the amount of $^{56}Ni$ produced (\\citealt{hmk93,khokhlov93,hoflich96b,mazzali01}). Thus, within the delayed detonation scenario, both normal and subluminous SNe~Ia can be explained as variants of a single phenomenon (H95, \\citealt{hkw95}, \\citealt{umeda99}). This makes the DD model an attractive scenario and it is used as the basis for our present analysis. In our models, the explosions LCs and spectra are all calculated self-consistently. Given the initial structure of the progenitor and a description of the nuclear burning front, the light curves and spectra are calculated directly from the explosion model without any additional parameters. Since the predicted observables follow directly from the supernova model, this approach provides a direct link between the observations and the explosion physics and progenitor properties. Finally, we note that detailed analyses of observed SN~Ia spectra and light curves indicate that mergers \\citep{benz90} and pure deflagration SNe (such as the W7 model) could contribute to the population of bright SNe~Ia (\\citealt{hk96}, hereafter HK96; \\citealt{hatano00}). On the other hand, it is not obvious that subluminous SNe~Ia and the brightness decline relation can be understood within pure deflagration models. In particular, the amount of burning, and consequently the explosion energy, will be strongly correlated with $^{56}Ni$ production. Since no detailed studies for deflagration models have been performed, these models cannot be ruled out but, as we discuss in \\S 5, the IR spectra of SN~1999by are hard to reconcile with a pure deflagration. \\subsection{Numerical Methods} \\subsubsection{The Explosion} Despite recent progress in our understanding of nuclear burning fronts, current 3-D models are not sufficiently evolved to allow for a consistent treatment of the burning front throughout all phases \\citep[see also \\S 4.3 \\& \\S 5]{khokhlov01}. Thus detailed models rely on parameterized descriptions guided by 3-D results. Within spherical models, the transition from deflagration to detonation can be conveniently parameterized by a density $\\rho_{tr}$ which has been found to be the dominating factor for the determining chemical structure, light curves and spectra. We have calculated explosion models using a one-dimensional radiation-hydro code (HK96) that solves the hydrodynamical equations explicitly by the piecewise parabolic method \\citep{collela84}. Nuclear burning is taken into account using an extended network of 606 isotopes from n,p to $^{74}Kr$ \\citep[and references therein]{thieleman96}. The propagation of the nuclear burning front is given by the velocity of sound behind the burning front in the case of a detonation wave, and in a parameterized form during the deflagration phase calibrated by detailed 3-D calculations \\citep[e.g.][]{khokhlov01}. We use the parameterization as described in \\citet{dominguez00}. For a deflagration front at distance $r_{burn}$ from the center, we assume that the burning velocity is given by $v_{\\rm burn}=max(v_{t}, v_{l})$, where $v_{l}$ and $v_{t}$ are the laminar and turbulent velocities with $$ v_{t}= 0.2 ~\\sqrt{\\alpha_{T} ~ g~L_f},~ \\eqno{[1]} $$ with~ $$\\alpha_T ={(\\alpha-1)/( \\alpha +1})$$ ~and ~ $$\\alpha ={\\rho^+(r_{\\rm burn})/ \\rho^-(r_{\\rm burn})}.$$ \\noindent Here $\\alpha _T$ is the Atwood number, $L_f$ is the characteristic length scale, and $\\rho^+$ and $\\rho^-$ are the densities in front of and behind the burning front, respectively. The quantities $\\alpha$ and $L_f$ are directly taken from the hydrodynamical model at the location of the burning front and we take $L_f=r_{\\rm burn(t)}$. The transition density is treated as a free parameter. \\subsubsection{The Light Curves} From these explosion models, the subsequent expansion, bolometric and broad band light curves (LC) are calculated following the method described by \\citet{hoflich98}, and references therein. The LC-code is the same used for the explosion except that $\\gamma$-ray transport is included via a Monte Carlo scheme and nuclear burning is neglected. In order to allow a more consistent treatment of the expansion, we solve the time-dependent, frequency-averaged radiation moment equations. The frequency-averaged variable Eddington factors and mean opacities are calculated by solving the frequency-dependent transport equations. About one thousand frequencies (in one hundred frequency groups) and about nine hundred depth points are used. At each time step, we use T(r) to determine the Eddington factors and mean opacities by solving the frequency-dependent radiation transport equation in a co-moving frame and integrate to obtain the frequency-averaged quantities. The averaged opacities have been calculated assuming local thermodynamic equilibrium (LTE). Both the monochromatic and mean opacities are calculated in the narrow line limit. Scattering, photon redistribution, and thermalization terms used in the light curve opacity calculations are taken into account. In previous works (e.g. \\citet{hk96}), the photon redistribution and thermalization terms have been calibrated for a sample of spectra using the formalism of the equivalent two level approach (H95). For increased consistency, we use the same equations and atomic models but solve the rate equations simultaneously with the light curves calculation at about every 100$^{th}$ time step, at the expense of some simplifications in the NLTE-part compared to H95. For the opacities we use the narrow line limit, and for the radiation fields we use the solution of the monochromatic radiation transport using $\\approx 1000$ frequency groups. \\subsubsection {Spectral Calculations} Our non-LTE code (H95, and references therein) solves the relativistic radiation transport equations in a co-moving frame. The spectra are computed for various epochs using the chemical, density, and luminosity structure and $\\gamma$-ray deposition resulting from the light curve coder. This provides a tight coupling between the explosion model and the radiative transfer. The effects of instantaneous energy deposition by $\\gamma$-rays, the stored energy (in the thermal bath and in ionization) and the energy loss due to the adiabatic expansion are taken into account. Bound-bound, bound-free and free-free opacities are included in the radiation transport which has been discretized by about $2 \\times 10^4$ frequencies and 97 radial points. The radiation transport equations are solved consistently with the statistical equations and ionization due to $\\gamma$-radiation for the most important elements and ions. Typically, between 27 to 137 bound levels are taken into account. We use \\ion{C}{1} (27/123/242), \\ion{O}{1} (43/129/431), \\ion{Mg}{2} (20/60/153), \\ion{Si}{2} (35/212/506), \\ion{Ca}{2} (41/195/742), \\ion{Ti}{2} (62/75/592), \\ion{Fe}{2} (137/3120/7293), \\ion{Co}{2} (84/1355/5396), \\ion{Ni}{2} (71/865/3064) where the first, second and third numbers in brackets denote the number of levels, bound-bound transitions, and number of discrete lines for the radiation transport. The latter number is larger because nearby levels within multiplets have been merged for the rates. The neighboring ionization stages have been approximated by simplified atomic models restricted to a few NLTE levels + LTE levels. The energy levels and cross sections of bound-bound transitions are taken from \\citet{kurucz93} starting at the ground state. The bound-free cross sections are taken from TOPBASE 0.5 as implemented in Stra\\ss burg \\citep{mendoza93}. Collisional transitions are treated in the `classical' hydrogen-like approximation \\citep{mihalas78} that relates the radiative to the collisional gf-values. All form factors are set to 1. About 10$^{6}$ additional lines are included (out of a line list of $4 \\times 10^7$) assuming LTE-level populations. The scattering, photon redistribution, and thermalization terms are computed with an equivalent-two-level formalism that is calibrated using NLTE models. \\subsection {Modeling Results} Spherical dynamical explosions, light curves, and spectra are calculated for both normal bright and subluminous SNe~Ia. Due to the one-dimensional nature of the models, the moment of the transition to a detonation is a free parameter. The moment of deflagration-to-detonation transition (DDT) is conveniently parameterized by introducing the transition density, $\\rho_{\\rm tr}$, at which DDT happens. Within the DD scenario, the free model parameters are: 1) The chemical structure of the exploding WD, 2) Its central density $\\rho_c$ at the time of the explosion, 3) The description of the deflagration front, and 4) The density $\\rho_{tr}$ at which the transition from deflagration to detonation occurs. Note that the density structure of a WD only weakly depends on the temperature because it is highly degenerate. The central density $\\rho_c$ depends mainly on the history of accretion. Model parameters have been chosen for the progenitor and $\\rho_c$ which allow us to reproduce light curves and spectra of `typical' Type Ia supernovae. In all models, the structure of the exploding C/O white dwarf is based on a model star with 5 M$_{\\sun}$ at the main sequence and solar metallicity. Through accretion, this core has been grown close to the Chandrasekhar limit (see Model 5p0y23z22 in \\citealt{dominguez01}). At the time of the explosion of the WD, its central density is 2.0$\\times 10^9$ g~cm$^{-3}$ and its mass is close to 1.37$M_\\odot$. The transition density $\\rho_{tr}$ has been identified as the main factor which determines the $^{56}Ni$ production and, thus the brightness of a SNe~Ia (H95; \\citealt{hkw95}; \\citealt{iwamoto99}). The transition density $\\rho_{tr}$ from deflagration to detonation is varied between 8 and 27 $\\times 10^6$ g~cm$^{-3}$ to span the entire range of brightness in SNe~Ia. The models are identified by {\\sl 5p0z22.ext} where {\\sl ext} is the value of $\\rho_{tr}$ in $10^6$ g~cm$^{-3}$. Some of the basic quantities are given in Table~\\ref{modelprop}. \\subsubsection{Explosion Models} Here we will restrict our discussion to the basic features of DD models which are relevant for the spectral analysis of SN~1999by. For a more detailed discussion of the hydrodynamical evolution of DD-models, see \\citet{khokhlov91}, H95, and references therein. During the deflagration phase, the WD is lifted, starts to expand, and after burning about 0.22 to 0.28 $M_\\odot$, the transition to a detonation is triggered. Figures \\ref{density} and \\ref{abundance} show the density, velocity and chemical structures for representative models after a homologous expansion has been achieved. The final velocity and density structures are rather similar. The expansion velocities decrease slightly with $\\rho_{tr}$ because a significant amount of oxygen remains unburned and did not contribute to the energy production. In 5p0z22.8, oxygen remained unburned in the outer $\\approx 0.4 M_\\sun $. Overall, the density is decreasing with radius because burning took place throughout the entire WD except in the very outer layers. In contrast, pure deflagration model such as W7 \\citep{nomoto84}, pulsating delayed detonation \\citep{hkw95}, and merger models \\citep{khokhlov93} show a significant density bump at the boundary between burned and unburned layers. Iron group elements are produced in the inner layers where density and temperature stay high for a sufficient time during the explosion, whereas intermediate mass elements are produced in the layers above where nuclear statistical equilibrium (NSE) has no time to set in. In the outer zones, products of explosive oxygen (e.g. Si, S) and explosive carbon burning (O, Mg, Ne) and some Si are seen. Only a very thin layer of unburned C/O remains. For layers with final velocities $v\\leq 3000$~km~s$^{-1}$, the densities in the early stages of the explosion are sufficiently high for electron capture. This results in the production of $^{54,56}$Fe \\& $^{58}$Ni whereas in the outer layers only $^{56}$Ni is produced. We employ the same description for the deflagration front in all models and, consequently, the very inner structures are almost identical in all models. A small dip in the Ni abundance is present at the DDT because the time scales for burning to NSE are comparable to the expansion time scales at densities of $\\approx 10^7$ g~cm$^{-3}$. This feature can be expected to be smeared out in multi-dimensional calculations \\citep{livne99}. After the DDT, a detonation front develops and the burning of fuel is triggered by compression of unburned material. Consequently, burning takes place under higher densities compared to a deflagration. In general, additional material is burned up to NSE and this extends the layers of $^{56}$Ni for all $\\rho_{tr} \\ge 1.2 \\times 10^7$ g~cm$^{-3}$. For models with smaller $\\rho_{tr}$, burning only proceeds up to Si despite the compression. This increases the production of intermediate mass elements (e.g. Si, S) at the expense of $^{56}$Ni. In our models, the efficiency for $^{56}$Ni production drops rapidly as a function of decreasing $\\rho_{tr}$ (see Table~\\ref{modelprop}). We note that for $\\rho_{tr} \\leq 6 \\times 10^{6}$ g~cm$^{-3}$, nuclear burning stops all together for $\\rho < \\rho_{th}$ and, unlike any observed SN, one would see a very slowly expanding C/O envelope heated by little $^{56}$Ni. There exists a rather strict lower limit for the absolute brightness of SNe~Ia within the Chandrasekhar mass models. This can be understood as follows. The absolute brightness of a SNe~Ia is mainly determined by the $^{56}Ni$. In general, $^{56}Ni$ is produced both during the deflagration and the detonation phase. For the most subluminous SNe~Ia however, no $^{56}Ni$ is produced during the detonation. Within the $M_{Ch}$ scenario and DD-models in particular, we need to burn a minimum of $\\approx 0.2 M_\\odot$ during the deflagration to achieve the required `lift/pre-expansion' of the WD. During the deflagration phase, iron group elements are produced. The fraction of $^{56}Ni$ depends on the amount of material which undergoes electron capture, i.e. on the central density and, to a smaller extent, on the description of the burning front (\\citealt{brachwitz01}; \\citealt{dominguez01}). For a model with $\\rho_c = 2 \\times 10^9$ g~cm$^{-3}$, we find a lower limit of $M(^{56}Ni)\\approx 0.09 M_\\odot$. Recently, \\citet{brachwitz01} studied the influence of $\\rho_{c}$ between 1.5 to $8 \\times 10^9$ g~cm$^{-3}$ on the nuclear burning for similar DD-models with a high transition density which may describe `normal' SN~Ia. In those models, the $^{56}Ni$ produced during the deflagration phase varies by $\\leq 25 \\%$ which provides an estimate for the possible variation in the minimum $^{56}Ni $ production. We conclude that there exists a lower limit of $M(^{56}Ni) \\approx 0.1 M_{\\sun} \\pm 25 \\%$ within DD-models. In principle, 3-D effects may change the lower limit for the $^{56}Ni$ production. However, all current 3-D calculations show that rising blobs are formed at high densities which undergo complete Si-burning because the burning time scales to Ni shorter than the hydro time scales. In the case of a DDT or a transition to a phase of very fast deflagration, the amount of burning will likely be similar to the 1-D case. Thus, we do not expect a significant change in the minimum ${56}Ni$ production. \\subsubsection{General properties of theoretical light curves} Optical light curves and spectra of SN~Ia are not the subject of this paper but we used their general properties for selecting a model for SN1999by. Within the DD scenario, both normal and subluminous SNe~Ia can be produced (H95, \\citealt{hkw95}, HKW95 hereafter; \\citealt{umeda99}). The amount of $^{56}$Ni ($M_{^{56}Ni}$) depends primarily on $\\rho_{tr}$ (H95; HKW95; \\citealt{umeda99}) and to a much lesser extent on the assumed value of the deflagration speed, initial central density of the WD, and initial chemical composition (ratio of carbon to oxygen). Models with smaller transition density give less nickel and hence both lower peak luminosity and lower temperatures (HKW95, \\citealt{umeda99}). In DDs, almost the entire WD is burned and the total production of nuclear energy is almost constant. This and the dominance of $\\rho_{tr}$ for determining the $^{56}Ni$ production are the basis of why, to first approximation, SNe~Ia appear to be a one-parameter family. The observed $M_V(\\Delta M_{\\Delta t=15d})$ relation can be understood as an opacity effect \\citep{hoflich96b}, namely, the consequence of rapidly dropping opacities at low temperatures \\citep{hmk93,khokhlov93,mazzali01}. Less Ni means lower temperature and consequently, reduced mean opacities because the emissivity is shifted from the UV toward longer wavelengths with less line blocking. A more rapidly decreasing photosphere causes a faster release of the stored energy and, as a consequence, steeper declining LCs with decreasing brightness. The DD models give a natural and physically well-motivated origin of the $M_V(\\Delta M_{\\Delta t=15d})$ relation of SNe~Ia (\\citealt{hamuy95,hamuy96b,phillips87,suntzeff99}). The broad band \\textit{B} and \\textit{V} light curves are shown in Figure~\\ref{lcs}. Some general properties are given in Table~\\ref{modelprop} and Figure~\\ref{dm15}. We note that the decline of the \\textit{B} and \\textit{V} LC during the early $^{56}Co$ tail is slightly smaller for the most subluminous models (e.g. 5p0z22.08) compared to the intermediate models (e.g. 5p0z22.16) because, at day 40, the $\\gamma $-rays are almost fully trapped in the former whereas, for the latter, the escape probability of $\\gamma $-rays is already increasing (HKW95). With varying $\\rho_{tr}$, we find rise times between 13.9 to 19.0 days, and the maximum brightness $M_V$ spans a range $-17.21^m$ and $-19.35^m$ with $B-V$ between $+0.47^m$ to $-0.03^m$. The value of $M_V$ is primarily determined by the transition density. Variations in the progenitors (main sequence mass and metallicity) or the central density at the point of ignition causes an spread around $M_V$ of $\\approx0.2~...~0.3^m$ \\citep{hoflich98,dominguez01}. As discussed above, the brightness decline relation $M_V(\\Delta M_{\\Delta t=15d})$ can be understood as an opacity effect or, more precisely, it is due to the drop of the opacity at low temperatures. Therefore, subluminous models are redder and the decline is steeper. In our series of models, $M_V(\\Delta M_{\\Delta t=15d})$ ranges between $1.45^m$ and $0.91^m$. However, $M_V(\\Delta M_{\\Delta t=15d})$ is not a strictly linear relation (Fig. \\ref{dm15}). For bright SNe~Ia ($-19.35^m \\leq M_V \\leq -18.72^m$), called {\\sl Branch-normals} (\\citep{branch96,branch99}), the linear decline is relatively flat, followed by a steeper decline toward the subluminous models. For normal-luminosity models, the photospheric temperature is well above the critical value at which the opacity drops ($\\approx 7000 ~... 8000 K$; \\citealt{hmk93}) whereas the opacity drops rapidly in subluminous models soon after maximum light. ", + "conclusions": "IR-spectra of the subluminous SN1999by have been presented which cover the time evolution from about 4 days before to 2 weeks before maximum light. This is the first subluminous SN Ia (and arguably the first SN~Ia) for which IR spectra with good time coverage are available. These observations allowed us to determine the chemical structure of the SN envelope. Based on a delayed detonation model, a self-consistent set of hydrodynamic explosions, light curves, and synthetic spectra have been calculated. This analysis has only two free parameters: the initial structure of the progenitor and the description of the nuclear burning front. The light curves and spectral evolution follow directly from the explosion model without any further tuning, thus providing a tight link between the model physics and the predicted observables. By varying a single parameter, the transition density at which detonation occurs, a set of models has been constructed which spans the observed brightness variation of Type~Ia supernovae. The absolute maximum brightness depends primarily on the $^{56}Ni$ production which, for DD-models, depends mainly on the transition density $\\rho_{tr}$. The brightness-decline relation $M_V (\\Delta M_{\\Delta t =15d})$ observed in SNe~Ia is also reproduced in these models. In the DD scenario this relation is a result of the temperature dependence of the opacity, or more precisely, as a consequence of the rapid drop in the opacities for temperatures less than about 7000 to 8000 K. Within $M_{Ch}$ WD models, a certain amount of burning during the deflagration phase is needed to pre-expand the WD and avoid burning the entire star to $^{56}Ni$. This implies a lower limit for the $^{56}Ni$ mass of about $0.1 M_\\odot$ $\\pm 25 \\%$ and, consequently, implies a minimum brightness for SNe~Ia within this scenario. The best model for SN~1999by ($\\rho_{tr} = 8 \\times 10^7$~g~cm$^{-3}$, selected by matching to the predicted and observed optical light curves) is close to this minimum Ni yield. The data are consistent with little or no interstellar reddening ($E(B-V) \\leq 0.12^m$), and the derived distance is $11 \\pm 2.5$~Mpc or $12 \\pm 1$~Mpc if we take the limit for $E(B-V)$ from the models or assume $E(B-V)=0.015^m$ according \\citet{schlegel98} for the galaxy. Without any further modification, this subluminous model has been used to analyse the IR-spectra from May 6, May 10, May 16 and May 24, 1999, which correspond to $-4$~d, 0~d, +7~d and +14~d after maximum light. The photosphere ($\\tau_{Thomson} = 1$) recedes from 14000, to 10500, 6500 and 4000 km~s$^{-1}$. The observed and theoretical spectra agree reasonably well with respect to the Doppler shift of lines and all strong features could be identified. Before maximum light, the spectra are dominated by products of about explosive carbon burning (O, Mg), and Si. Spectra taken at +7~d and +14~d after maximum are dominated by products of incomplete Si burning. At about 2 weeks after maximum, the iron-group elements begin to show up. The long duration of the phases dominated by layers of explosive carbon burning and incomplete Si burning implies massive layers of these burning stages that are comparable with our model results. ($\\approx 0.45 $ and $ 0.65 M_\\odot $ for explosive carbon and incomplete silicon burning, respectively.) This, together with the $^{56}Ni$ mass, argues that SN~1999by was the explosion of a WD at or near the Chandrasekhar mass. Finally, we note that the observed IR spectra are at odds with recent 3-D calculations \\citep{khokhlov01} which predict that large scale chemical inhomogeneities filling 50 -- 70\\% (in mass) of the WD will be formed during the deflagration phase. When the effect of such inhomogeneous mixing is tested by mixing the inner layers of our SN~1999by model, significant differences appear between the model spectra and the observed data. This suggests that no significant large-scale mixing took place in SN~1999by. The lack of observed mixing and the asphericity seen in SN~1999by may be important clues into the nature of subluminous SNe~Ia, and may be related to the reason for a low DDT transition density (and hence the low luminosity). Alternatively, we may have an extended smoldering phase of the WD prior to the explosion which skips the deflagration phase altogether (see \\S 5). In either case, the pre-conditioning of the WD prior to the explosion seems to be a key for understanding SNe~Ia." + }, + "0112/math0112296_arXiv.txt": { + "abstract": "Two points are randomly selected inside a three-dimensional euclidian cube. The value $l$ of their separation lies somewhere between zero and the length of a diagonal of the cube. The probability density ${\\cal P}(l)$ of the separation is constructed analytically. Also a Monte Carlo computer simulation is performed, showing good agreement with the formulas obtained. ", + "introduction": "An important problem in geometry and statistics is: given a convex compact space endowed with a metric, and randomly choosing two points in the space, find the probability density $\\Pl$ that these points have a specified separation $l$. The study of this problem has a long history \\cite{Beres}, and recently gained considerable impetus from researchers in cosmic crystallography \\cite{FarrarMelott}-\\cite{Janna}. In a recent paper the functions $\\Pl$ corresponding to $2D$ disks and rectangles were obtained \\cite{Memb}. The methodology introduced in that work is here extended to a $3D$ euclidian cube. ", + "conclusions": "" + }, + "0112/astro-ph0112304_arXiv.txt": { + "abstract": "V651 Mon is the binary central star of the bipolar planetary nebula NGC 2346. The star showed the second-ever deep fading in 1996--1997, which was presumably caused by obscuration by a dust cloud in the planetary nebula, as was proposed to explain the 1981--1985 event. The entire duration of the 1996--1997 event was $\\sim$400 d, remarkably shorter than the 1981--1985 event, suggesting that the obscuring body was smaller or had a larger tangential velocity. The most remarkable feature in this event was the presence of a sharply defined transient clearing (brightening). From the time-scale of the variation, we propose an upper limit of the projected scale of several times $\\sim$10$^{11}$ cm of the structure responsible for the brightening. This observation provides the first evidence for a sharply defined, small lucent structure within the obscuring body around the central star of NGC 2346 ", + "introduction": "V651 Mon is the binary central star of the bipolar planetary nebula NGC 2346. The 15.991-d binary consists of an A5 V star, which emits most of the visual light, and a hot ($T_{\\rm eff}$ $\\sim$10$^5$ K) subdwarf, which emits the most of the ultraviolet light \\citep{men81}. This binary has received special attention since its fading episode starting in 1981 \\citep{koh82}. The fading initially showed a strong modulation with the orbital period. Subsequent observations led to a picture of a passing cloud (within the planetary nebula) in front of the binary: eclipse-like fadings in the visual light occur when the orbiting A5 V star passes behind the cloud (for a review on this object, see \\cite{cos86}). \\citet{cos86} expect, from the amount of dust from IRAS observations, that the probability of such passages is very small. In fact, plate searches could not detect any similar events between 1899 and 1981 \\citep{sch83}. Further detailed photoelectric observations only caught small variations \\citep{koh92}, or a transient short fading \\citep{gro93}. In spite of a prediction by \\citet{cos86}, the second-ever deep fading was reported in 1996 September \\citep{ove96}. We started CCD observations upon this alert. ", + "conclusions": "Figure \\ref{fig:fading} presents the overall course of the 1996--1997 fading episode. Visual observations were reported to VSNET\\footnote{$\\langle$ http://www.kusastro.kyoto-u.ac.jp/vsnet/ $\\rangle$.} (VSNET is an international network of variable star observers, designed to make collaborative studies on selected variable stars). All observers used comparison stars calibrated in the $V$-band. The present CCD $V$-band observations are also shown (open circles). The total duration of the event was $\\sim$400 d, which is remarkably shorter than the 1981--1985 event. Figure \\ref{fig:lc} shows an overall light curve from our CCD observations. The CCD observation started when the 1996--1997 fading episode reached close to its minimum. The object recovered to its normal magnitude during the latest two observations (at the end of 1997). \\begin{figure*} \\begin{center} \\FigureFile(130mm,90mm){fig1.eps} \\end{center} \\caption{Visual light curve of the 1996--1997 fading from observations reported to VSNET (filled squares). CCD $V$-band observations (open circles) from table \\ref{tab:log} are also shown. The total duration of the deep fading episode was $\\sim$400 d.} \\label{fig:fading} \\end{figure*} \\begin{figure} \\begin{center} \\FigureFile(88mm,60mm){fig2.eps} \\end{center} \\caption{Light curve of V651 Mon from the present CCD observations.} \\label{fig:lc} \\end{figure} Following the interpretation by \\citet{cos86}, the cloud responsible for the fade was either smaller or had a larger tangential velocity. The total depth of the fading ($\\sim$2.7 mag) may be slightly smaller than the extreme value reported in the 1981--1985 event. This difference, however, may be a result of a systematic difference of the measurements: in the 1981--1985 event, the aperture background subtraction technique in photoelectric photometry was employed in order to measure the faint object against the bright nebular background, while the present observations of the 1996--1997 event used PSF fitting of the stellar profile on two-dimensional CCD images. By taking this effect into account, the depth was found to be roughly comparable to that of the 1981--1985 event, implying that the opacity of the cloud is similar. Another notable feature of the present event is the presence of a sharply defined transient brightening around JD 2450540. The enlarged light curve around this brightening is shown in figure \\ref{fig:brightening}. Although visual observations reported to VSNET suggest the presence of a weak 16-d periodicity, as was observed in the 1981--1985 event, no other brightening with a similar amplitude was observed around the binary phase of the brightening. Since the binary motion of the central close binary is expected to be much faster than the tangential velocity of the orbiting cloud [\\citet{cos86} proposed 0.14 km s$^{-1}$ for the 1981--1985 event], the sharply defined brightening is considered to more reflect the orbital motion of the A5 V star against the foreground cloud. By assuming $K_1$=16.4 km s$^{-1}$ \\citep{men81}, the sharp structure of the light curve corresponds to an upper limit of the projected scale of several times $\\sim$10$^{11}$ cm. This scale is by a factor of about ten smaller than the entire cloud responsible for the 1981--1985 fade, which was estimated to have a size of 2--5 $\\times$ 10$^{12}$ cm \\citep{cos86}. This observation not only provides the first evidence for such a sharply defined, small lucent structure within the obscuring body around the central star of NGC 2346, but also the first-ever direct determination of the dimension of a fine structure within a dust cloud of a planetary nebula. \\begin{figure} \\begin{center} \\FigureFile(88mm,60mm){fig3.eps} \\end{center} \\caption{Transient, sharp brightening observed during the fading episode.} \\label{fig:brightening} \\end{figure} \\vskip 3mm The authors are grateful to VSNET members for providing vital observations." + }, + "0112/astro-ph0112132_arXiv.txt": { + "abstract": "Central cool gas component that is often observed from a well-relaxed cluster system has long been interpreted as a consequence of ``Cooling Flow'' (CF), radiative cooling followed by inflow of Intra-Cluster Medium (ICM). However, recent XMM-Newton spectroscopy has shown no signatures of cooler gas phases below certain temperatures in typical CF clusters (A1795, Tamura et al. 2001; A1835, Peterson et al. 2001). This contradicts the conventional CF model or at least requires a major revision of the model. In order to investigate statistical properties of the central cool component, we performed systematic analysis of {\\it ASCA} data on 85 clusters. We found that 1) temperature of the central cool component strongly depends on the temperature of the main ICM, 2) the cool component is selectively found around a brightest cluster galaxy (BCG) that coincide with the X-ray peak position, and 3) the luminosity-temperature ($L$--$T$) relation of the cool component shows nice agreement with the $L$--$T$ relation of the main ICM. Together with the previous observational fact that, in some of the ``CF'' clusters, the total gravitating mass is clustering in two distinct spatial scales, a main cluster component and a second small-scale system, we conclude that the central cool component is associated with the second small-scale self-gravitating system that is immersed in the host cluster, and the cool component temperature reflects the gravitational potential depth. ", + "introduction": "For the study we used brightest 85 clusters observed with {\\it ASCA} whose redshifts distribute from 0.004 to 0.201 with a median of 0.046. From each GIS and SIS instrument, a spectrum over a central region of $2'$ radius and one outside $2'$ radius were accumulated. We fitted each set of {\\it ASCA} spectra with the two-temperature model (2T model), in which a hot component is filling entire cluster region, while in the central region a cool component is allowed to coexist with the hot component forming two-phase plasma. We then obtained X-ray luminosity in 0.1-2.4~keV band and temperature of cool and hot component, $L_{\\rm c}$, $T_{\\rm c}$, $L_{\\rm h}$, and $T_{\\rm h}$. Full results are given in Ikebe et al. (2001). An example of the 2T model fit is shown in Fig.~1 for Abell~1795, which has one of the brightest central cool component. \\begin{figure} \\centerline{ \\psfig{file=ikebe1fig1.ps,width=6.0cm,angle=-90,clip=} \\psfig{file=ikebe1fig2.ps,width=6.0cm,angle=-90,clip=}} \\caption{ {\\it ASCA} spectra of Abell~1795 extracted from the center (left) and outer (right) regions. Crosses and filled circles show SIS and GIS data, respectively. Flux mixing effects due to the largely extended point-spread functions of X-Ray Telescope have been taken into account and the four spectra are fitted simultaneously with the 2T model. The best-fit model for the SIS and GIS are illustrated with solid and dotted lines respectively.} \\end{figure} ", + "conclusions": "Based on the results from the 2T model fitting to the {\\it ASCA} data of the 85 selected clusters, we investigated various correlations. In the study, we classified the cluster sample into three groups, which are Strong Cool Component (SCC), X-ray Dominant (XD) and non-X-ray Dominant (nXD). An SCC cluster is defined as one showing a very strong (statistically very significant) cool component in the {\\it ASCA} spectrum. Among the rest of the clusters, one in which the X-ray peak coincides with the Brightest Cluster Galaxy (BCG) is classified as an XD cluster, while nXD is defined as one without such a galaxy. All the SCC clusters turned out to be XDs, too. \\begin{figure} \\centerline{ \\psfig{file=ikebe3fig1.ps,width=5.5cm,angle=-90,clip=} \\psfig{file=ikebe3fig2.ps,width=5.5cm,angle=-90,clip=}} \\caption{ (a) The $L_{\\rm c}$--$T_{\\rm c}$ relation for the SCC clusters. The line is the best-fit power-law function to the $L_{\\rm h}$--$T_{\\rm h}$ relation illustrated in Fig.~6. (b) The $L_{\\rm c}$--$T_{\\rm c}$ relation for all the sample clusters. Filled circles, open circles, and filled stars specify SCC, XD, and nXD clusters, respectively.} \\end{figure} \\begin{figure} \\centerline{\\psfig{file=ikebe4fig1.ps,width=5.5cm,angle=-90,clip=}} \\caption{$T_{\\rm c}$--$T_{\\rm h}$ relation for the SCC clusters. The line represents $T_{\\rm c}$=$T_{\\rm h}/2$.} \\end{figure} For the SCC clusters, we determined the temperatures and emission measures of each cool and hot component. Figure~3(a) shows the cool component luminosity as a function of the cool component temperature for the SCC clusters. A positive correlation is clearly seen. Figure~4 shows the correlation between the two temperatures of the hot and cool component for the SCC clusters. Surprisingly, the ratios between the two temperatures, $T_{\\rm c}/T_{\\rm h}$, is virtually constant and the relation, $T_{\\rm c}$=$T_{\\rm h}/2$ well represents the correlation. For clusters that do not show very strong cool component, i.e. non-SCC clusters, we fixed $T_{\\rm c}$ at the half value of each mean temperature derived with a single-temperature model fitting, according to the $T_{\\rm c}$=$T_{\\rm h}/2$ relation. The luminosity of the cool component thus estimated for the non-SCC clusters is illustrated in Fig.~3(b). The results of the SCC clusters are also overlayed. SCC and XD clusters exhibit systematically more luminous cool component than nXD clusters. Actually, many of nXD clusters give only upper limit to $L_{\\rm c}$. We obtained the $L-T$ relation of the hot component, which is illustrated in Fig.~5(a). Unlike for the cool component, the three classes do not show any systematic difference. It is then clearly shown that the X-ray characteristics of the XD and nXD clusters are segregated mainly by the central cool component. To compare with the $L-T$ relation of the cool component would be very interesting. In the Fig.~5(b), the $L_{\\rm c}$--$T_{\\rm c}$ relation for the SCC clusters shown in Fig.~3 is overlayed on the $L_{\\rm h}$--$T_{\\rm h}$ relation. A surprising agreement is clearly seen. The best fit power-law function to the $L_{\\rm h}$--$T_{\\rm h}$ relation is compared with the SCC $L_{\\rm c}$--$T_{\\rm c}$ relation in Fig.~3. The similarity of the two $L$--$T$ relation suggests that both the cool and hot components are related to individual gravitational bound objects. Therefore, the cool component can be naturally interpreted as ICM filling a self-gravitating system whose size is comparable to a giant elliptical galaxy or a group of galaxies, which is immersed in the host cluster. \\begin{figure} \\centerline{ \\psfig{file=ikebe5fig1.ps,width=5.5cm,angle=-90,clip=} \\psfig{file=ikebe5fig2.ps,width=5.5cm,angle=-90,clip=}} \\caption{(a) The $L_{\\rm h}$--$T_{\\rm h}$ relation for all the sample clusters. Filled circles, open circles, and filled stars specify SCC, XD, and nXD clusters, respectively. (b) The $L_{\\rm c}$--$T_{\\rm c}$ relation for the SCC clusters (filled circles) is overlayed on the $L_{\\rm h}$--$T_{\\rm h}$ relation (open circles).} \\end{figure}" + }, + "0112/astro-ph0112418_arXiv.txt": { + "abstract": "We study the possibility for an observational detection of oscillations due to baryons in the matter power spectrum and suggest a new cosmological test using the angular power spectrum of halos. The {\\it standard rulers} of the proposed test involve overall shape of the matter power spectrum and baryon oscillation peaks in projection, as a function of redshift. Since oscillations are erased at non-linear scales, traces at redshifts greater than 1 are generally preferred. Given the decrease in number density of clusters at high redshift, however, one is forced to use tracers corresponding to galaxy groups and galaxies themselves. ", + "introduction": "There is now increasing evidence for the presence of oscillations in the angular power spectrum of cosmic microwave background (CMB) anisotropies. The matter power spectrum of the large scale structure also contains a signature of baryons in the form of oscillations; the amplitudes of oscillations, however, are significantly lower with fluctuations at the level of $\\sim$ 5\\% for $\\Lambda$CDM with widths of order $\\Delta k \\sim 0.02$ h Mpc$^{-1}$. For an unambiguous detection, 3d surveys should have resolution scales greater than $L \\sim 2\\pi/\\Delta k \\sim$ 300 h$^{-1}$ Mpc in all three dimensions. In 2d, oscillations can also be detected via the projected power spectrum. The 2d power spectrum requires no precise redshifts and can be constructed with halos detected in upcoming wide-field lensing and SZ surveys. Following our suggestions that the angular power spectrum of halos can be used as a probe of cosmology (Cooray et al. 2001), projected oscillations provide a new cosmological test; this is similar to the measurement of the angular diameter distance to the last scattering surface using the first acoustic peak in the CMB anisotropy power spectrum. ", + "conclusions": "" + }, + "0112/astro-ph0112242_arXiv.txt": { + "abstract": "{The Crab-like SNR G21.5-0.9 was observed in the X--ray band (0.5-10 keV) by the {\\it XMM-Newton} satellite for over 100 ks. The large effective area of the {\\it EPIC} instrument has allowed us to to perform a deep search for pulsations from the central core of \\ee. No pulsations were found with upper limits on the pulsed fraction between 7.5 \\% and 40 \\% (depending on frequency and energy range). ", + "introduction": "Almost 5\\% of the $\\sim$ 225 known supernova remnants (SNRs) in our galaxy are classified as ``Crab-like'' or ``plerionic'' (Green 2000). From the spatial point of view, they are characterized by compact, centre-filled radio and X--ray morphology (Weiler \\& Panagia 1978). These SNRs show featureless power-law spectra, with a relatively flat spectral index in the radio regime ($\\alpha_{r}\\sim$0.0 - 0.3) and a steeper one at shorter wavelengths, which are typical of synchrotron processes. These morphological and spectral characteristics are explained by the presence of a central pulsar, which injects high energy electrons that suffer synchrotron radiation losses as they diffuse through the surrounding magnetic field (see, e.g., Reynolds \\& Chanan 1984). \\ee shows many characteristics of the Crab-like remnants. Both in the radio band (Morsi \\& Reich 1987) and at X--ray energies (Becker \\& Szymkowiak 1981) its emission is centrally peaked. Evidence for a non-thermal X--ray spectrum was also indicated by {\\it GINGA} observations (Asaoka \\& Koyama 1990). The available measurements of the neutral hydrogen absorption give a distance of $\\sim$ 4.8 kpc (Becker \\& Szymkowiak 1981). The radio luminosity of \\ee is $\\sim$ 1.8$\\times$10$^{34}$ {\\it d}$^{2}_{5}$ erg s$^{-1}$ (Morsi \\& Reich 1987), i.e. a factor $\\sim$ 9 smaller than that of the Crab, but its X--ray luminosity is a factor $\\sim$ 100 less, therefore the {\\it L$_{X}$/L$_{r}$} ratio is significantly lower. X--ray observations performed with the \\textit{Chandra} satellite detected a compact central core of $\\sim$ 2$''$ in size, at the center of the more extended synchrotron nebula of $\\sim$ 30$''$ radius (Slane et al. 2000). The central core, which is spatially resolved, most likely marks the position of the pulsar powering \\ee. Also a fainter, more extended ``halo'' (radius $\\sim 2'$) was detected with \\textit{Chandra}. This was tentatively interpreted as the outer ``shell'' formed by the expanding ejecta and the passage of the supernova-driven blastwave (Slane et al. 2000). However, a recent {\\it XMM-Newton} observation shows that also the halo has a non-thermal spectrum; it is probably a low surface brightness extension of the plerionic nebula (Warwick et al. 2001). However, this interpretation is not supported by radio data, since no significant radio emission has yet been detected at such a large distance from the source core (Bock \\& Wright 2001). Up to now no pulsed emission has been detected in the radio or X--ray energy range from the putative neutron star at the center of \\ee (Frail \\& Moffet 1993; Kaspi et al. 1996; Biggs \\& Lyne 1996; Slane et al. 2000). In this paper we report the results of a sensitive timing analysis on data provided by four {\\it XMM-Newton} observations. ", + "conclusions": "Our upper limits can be compared with those obtained by some recent works based on X--ray data of \\ee. Note that the values reported by Warwick et al. (2001), 3.5\\% and 5.5\\% (respectively for MOS+PN and PN only data), refer to the \\textit{total} flux within an extraction radius of 8$''$, including the nebular emission. Our results for the on-axis observation alone (the one used by these authors) are similar, but we reached a better sensitivity in the summed power spectra (corresponding to a factor $\\sim$4 greater exposure time). For instrumental reasons, our search was limited to periods longer than 146 ms. Although it is clearly possible that the pulsar in G21.5-0.9 has a shorter period, we note that Safi-Harb et al. (2001), from energetic considerations, estimated a period of P = 0.144 $(I_{45} / (\\dot E_{37} \\tau_{3}))^{1/2}$ s (where $I=10^{45}I_{45}$ g cm$^{2}$ is the moment of inertia, $\\dot E = 10^{37} \\dot E_{37}$ erg s$^{-1}$ is the spin-down energy loss and $\\tau = \\tau_{3}\\times$(3 kyr) is the pulsar characteristic age). For reasonable values of $\\tau$ and $\\dot E$, such a period falls in the range we could explore with EPIC. The same authors analyzed five data sets (total 75 ks) obtained with the \\textit{Chandra HRC} instrument. They report an upper limit of 16\\%, without quoting the confidence level and, presumably, referring to the total flux within 2$''$. The ``canonical'' picture of plerionic supernova remnants is based on young, energetic neutron stars with short rotation periods, such as the Crab pulsar (P=33 ms) or the recently discovered pulsar in 3C 58 (P=66 ms, Murray et al. 2001). However, other results show that also sensitive searches for slower pulsars are relevant: there are in fact relatively young pulsars with long periods. Besides the well known example of PSR B1509--58 in the SNR G320.4-01.2 (P=150 ms), other recent findings include the 325 ms pulsar in the SNR Kes 75 (Gotthelf et al. 2000), and PSR J1119--6127 (P=407 ms, Pivovaroff et al. 2001), which however does not have a bright synchrotron nebula. \\newpage" + }, + "0112/astro-ph0112074_arXiv.txt": { + "abstract": "As a logical next step in improving the radio-to-submm spectral index as a redshift indicator (Carilli \\& Yun), we have investigated a technique of using the entire radio-to-FIR spectral energy distribution (SED) for deriving photometric redshifts for dusty starburst galaxies at high redshift. A dusty starburst SED template is developed from theoretical understanding on various emission mechanisms related to massive star formation process, and the template parameters are selected by examining the observed properties of 23 IR selected starburst galaxies: $T_d=58$ K, $\\beta=1.35$, and $f_{nth}=1$. The major improvement in using this template SED for deriving photometric redshifts is the significant reduction in redshift uncertainty over the spectral index technique, particularly at higher redshifts. Intrinsic dispersion in the radio and FIR SEDs as well as absolute calibration and measurement errors contribute to the overall uncertainty of the technique. The derived photometric redshifts for five submm galaxies with known redshifts agree well with their spectroscopic redshifts within the estimated uncertainty. Photometric redshifts for seven submm galaxies without known spectroscopic redshifts (HDF850.1, CUDSS14.1, Lockman850.1, SMM~J00266+1708, SMM~J09429+4658, SMM~J14009+0252, FIRBACK~J1608+5418) are derived. ", + "introduction": "Sensitive observations at submm wavelengths are revealing what may be a population of active star forming galaxies at high redshift which are unseen in deep optical surveys due to dust obscuration \\citep{Smail97,Barger98,Hughes98,Eales99}. The differential source counts clearly indicate a large excess of far-infrared sources by a factor of 10-50 over the no-evolution models derived from the optical deep survey data \\citep{guiderdoni98,Blain99a,Matsuhara00,Scott01}, and a large fraction of star formation may be hidden by dust. The analysis of the submm SCUBA source counts and inferred redshift distribution based on the radio-submm flux density ratio \\citep{Carilli99} suggests that luminous dusty galaxies may dominate the star formation history at early epochs ($z\\sim 1-3$) \\citep{barger00}. Many of these faint submm sources are found to be very faint, red ($R\\ge25, K\\ge21$) sources \\citep{Smail99a}, whose optical redshifts may be inaccessible even for the 10-m class telescopes. These galaxies are nearly entirely missing from the optical studies of star formation at high redshifts \\citep [e.g.][]{Steidel99, Chapman00}, and a significant revision to the optically derived cosmic star formation history may be needed \\citep [see][]{Blain99a}. Because direct redshift measurements are not possible in most cases, the redshift distribution and the cosmic evolution of the dusty submm galaxy population are not well determined at the moment. In a recent paper, we have proposed a technique of using the radio-to-submm spectral index as a redshift indicator. This technique is based on the universal radio-to-far infrared (FIR) correlation for star forming galaxies \\citep{Condon92}, with the assumption that the spectral shapes may be similar enough to be able to differentiate between low and high redshift objects \\citep [see][]{Carilli99}. To understand the magnitude of scatter in this relation and resulting uncertainty in the redshift estimates, we derived an empirical $\\alpha^{350}_{1.4} - z$ relation\\footnote{$\\alpha^{350}_{1.4}$ is the spectral index between 1.4 GHz and 350 GHz.} using the observed spectral energy distributions (SEDs) of 17 low redshift starburst galaxies (Carilli \\& Yun 2000; also see Dunne, Clements, \\& Eales 2000a). There is a significant scatter, but the existing data on high redshift star forming galaxies and AGN-hosts appear to follow this relation well. While the radio-to-submm spectral index has been demonstrated to be a useful redshift indicator, using it to derive a redshift for any particular object may be risky because of the scatter in the observed $\\alpha^{350}_{1.4} - z$ relation and the flattening in the relation at high redshifts. Obtaining redshift estimates using flux density ratios at other wavelength bands has also been suggested although none is particularly more successful \\citep[e.g.][]{Hughes98,Blain99b,Hines99,Fox01}. More accurate redshift estimates for dusty starbursts may be obtained by utilizing the information associated with the {\\it entire radio-to-FIR SED}. Such a photometric redshift technique requires a good SED template, and here we derive an SED template based on the theoretical expectations of thermal dust, thermal Bremsstrahlung, and non-thermal synchrotron emission from a dusty starburst galaxy. Our SED template is constructed by examining the properties of the 23 IR selected starburst galaxies whose submm/FIR SED data are readily available in the literature. When applied to a sample of submm galaxies with known redshifts, the resulting redshift estimates are in excellent agreement with the spectroscopic redshifts. We also report photometric redshifts for a sample of submm galaxies without known spectroscopic redshifts, including the brightest SCUBA source in the Hubble Deep Field \\citep{Hughes98}. ", + "conclusions": "} The above discussions of the SED modeling for the individual submm detected galaxies clearly indicate that systematic SED variation among the individual galaxies is the dominant factor over the statistical uncertainty associated with the SED measurements. Here we discuss in detail the nature and the magnitude of the systematic effects associated with the dust and radio emission as well as the SED sampling and the intrinsic calibration uncertainties in the measurements. Their combined contribution to the photometric redshift uncertainty is quantified in \\S~\\ref{sec:err-total}. \\subsection{Effects of Dust Temperature and Emissivity \\label{sec:err-dust}} Doppler shift of a dust spectrum to a higher redshift is exactly equivalent to lowering of dust temperature, and departure in $T_d$ from the template SED translates directly to a redshift error \\citep[see][for further discussions]{Blain99b}. In other words, if the dust temperature of a galaxy is actually higher than our template, then the resulting $z_{ph}$ would be an under-estimate. While inherently subject to an object-to-object variation, the magnitude of uncertainty due to the spread in $T_d$ is well understood quantitatively: $\\Delta z \\sim {{\\Delta T_d}\\over{T_d}}(1+z)$. If $T_d$ is higher or lower by 10 K, the resulting error in the photometric redshift is about 0.3 at $z=1$, and it grows as $(1+z)$ reaching about 0.7 at $z=3$. Dust emissivity $\\beta$ is as important as dust temperature $T_d$ because it determines the location of the low frequency edge of the thermal dust feature in the SEDs. There is significant degeneracy between $\\beta$ and $T_d$ \\citep[e.g.][]{Dunne01}, and the decoupling the two becomes very difficult even when the dust spectrum is well sampled. In practice only a few measurements along the rising part of the dust SED are measured, and any deficiency in $T_d$ is at least partly offset by the compensating effect of $\\beta$. A trend supporting this effect is seen in Figure~\\ref{fig:betaTd} as a broad trend of decreasing $\\beta$ with increasing $T_d$. A better demonstration of this effect is seen in Figure~\\ref{fig:a220sed} where the submm SED points alone might be fit well by a dust spectrum with $T_d\\sim 45$ K and $\\beta\\sim 1.5$. The observed trend of increasing $T_d$ and deceasing $\\beta$ with increasing $SFR$ is also clearly seen in the theoretical modeling of dust heating and emission using a self-consistent 3-D radiative transfer code by \\citet[][]{Misselt01}. Therefore the error estimate based on the $T_d$ alone might serve only as a reasonable upper bound. \\subsection{Effects of the Radio Spectrum \\label{sec:err-radio}} Another reason why the photometric redshift error based on the scatter in $T_d$ alone might be an over-estimate is that our SED fitting scheme relies as much on the spectral trough between the radio continuum and thermal dust emission as the dust spectrum itself. If this spectral trough is considered as the primary SED feature from which the photometric redshift information is derived, then it is easy to see that the variation in the dust spectrum is strongly modulated by the radio spectrum which is not affected. The radio spectrum has its own uncertainty in its overall scaling as shown by the distribution of $f_{nth}$ (see Figure~\\ref{fig:fnth}), and its contribution to the photometric redshift error has already been alluded to in \\S~\\ref{sec:tests}. Presence of a powerful radio AGN was a severe limiting factor for our earlier redshift estimation based on the radio-to-submm spectral index \\citep{Carilli99}, and it is not surprising that radio AGNs pose the biggest obstacle for the SED fitting photometric technique as well. The photometric redshifts derived for an ensemble of QSOs that are associated with dusty hosts are given at the bottom of Table~\\ref{tab:ztable} (also see Fig.~\\ref{fig:comparez}). Although there are exceptions such as LBQS~1230+1627 where $z_{ph} \\sim z_{sp}$, the new photometric method does not fare much better than the spectral index method -- the median $\\chi^2_n \\sim 7$. In practice, radio AGNs are not likely to limit the photometric redshift technique utilizing a starburst SED template because only a small minority of the known submm galaxies show evidence of hosting a powerful AGN. Comparisons of deep X-ray and submm surveys have shown very little overlap between the detected sources \\citep[e.g.][]{Fabian00,Horns00,Barger01}. At least three out of five submm galaxies discussed in \\S~\\ref{sec:tests} are known to be an AGN host, but the presence of an AGN also makes their spectroscopic redshift measurement possible and the use of a photometric technique unnecessary. The total fraction of infrared selected galaxies whose radio-to-FIR SED shows a clear sign of energetic AGN is only a few percent in the local universe \\citep{Yun01}, and this fraction appears to remain roughly the same at higher redshifts. \\subsection{SED Sampling and Calibration \\label{sec:err-cal}} Another important source of uncertainty is poor sampling of the SED. Sampling only the radio part of the SED offers virtually no redshift information, and the same is true if only the submm part of the SED is measured. On the other hand, because we are fitting the SED features changing over logarithmic scales, only a few well placed SED data points are needed to derived the redshifts \\citep{Carilli99}. More than one measurement on both side of the SED trough is highly desirable for a successful photometric redshift determination. A related problem is the calibration and the relative weighting of individual SED measurements. Measurement uncertainties found in the literature often do not fully account for the overall calibration accuracy as demonstrated by the magnitude of the scatter and the sizes of the error bars plotted for the SED measurements of Arp~220 in Figure~\\ref{fig:a220sed}. For example, the IRAS 60 $\\mu$m and 100 $\\mu$m measurements found in the Point Source Catalog (PSC) and the Faint Source Catalog (FSC) are known to have systemic calibration differences of order 10\\% while the flux densities reported in the FSC often include uncertainties of only a few percent. For deriving the SED template from the 23 infrared selected starbursts, we adopted flux densities in the FSC but increased measurement uncertainties to 10\\% (unless the reported uncertainty was larger). This re-calibration of the measurement uncertainty is more than cosmetic since the relative weight of the data points are directly reflected in the $\\chi^2$ minimization. We found this re-weighting to be critically important in utilizing the submm measurements. Similarly, all radio continuum measurements are assigned at least 10\\% uncertainty in order to account for the overall uncertainty in the flux calibration. No effort was made to re-calibrate any of the submm data points since most measurements generally carry relatively large fractional uncertainties, but some are clearly under-estimates based on the SED plots such as shown in Figure~\\ref{fig:a220sed}. The same problems also plague the reported SED measurements of many submm galaxies, and they will inevitably impact the accuracy of the photometric redshifts derived from these data. One way to improve the situation in the future would be employing a frequency selective bolometer \\citep[e.g. ][]{Kowitt96,Meyer01} that can make simultaneous measurements of several submm bands with accurate relative calibration between the measurement bands. \\subsection{Estimate of Uncertainty in the Radio-to-FIR SED Technique \\label{sec:err-total}} The uncertainty of $\\sigma_{z}\\sim 0.3(1+z)$ derived from the scatter in dust temperature is probably the upper bound to the error one may expect from the photometric technique using the dusty starburst SED template. The compensating effects of the dust emissivity and the radio continuum are more difficult to quantify as are the uncertainties contributed by the scatter in $f_{nth}$. One way to estimate the collective uncertainty of the SED photometric technique is to apply this technique to the same 23 dusty starbursts from which the SED template was derived. This is not entirely circuitous since our SED template is based only on the average properties of these galaxies and has no knowledge of dispersions. At the least, we may be able to confirm the impact of the 10 K scatter in the dust temperature on the dispersion in $z_{ph}$ if dust temperature variation dominates the uncertainty in determining photometric redshifts. The resulting ``photometric redshift'' $z_{ph}$ for the 23 dusty starbursts are listed in the last column of Table~\\ref{tab:23gals}. There are several galaxies with negative $z_{ph}$ since the impact of higher dust temperature for a $z=0$ galaxy with $T_d > 58$ K would result in $z_{ph}<0$. The $\\chi^2$ minimization program is modified to search a redshift range of $-1\\ge z_{ph} \\ge +1$ with modified scaling along the flux density axis in order to remove the non-physical impact of negative redshift on $D_L$. A handful of galaxies with negative redshift are known (e.g. M81), but negative redshift is generally considered non-physical. Here, the only physically meaningful interpretation of a negative $z_{ph}$ we derive is the measure of the magnitude of departure in the radio and dust property from the template SED, just the same way any uncertainty in $z_{ph}$ should be interpreted at any redshift. A histogram of the resulting photometric redshifts shown in Figure~\\ref{fig:z0histo} suggests that the median ``redshift'' is about +0.05, suggesting the template SED may be slightly biased, but this offset is probably not very significant given the range of ``redshifts'' derived. ``Photometric redshifts'' as large as +0.3 and as small as $-$0.3 are found as expected from the scatter in $T_d$. The two extreme negative ``$z_{ph}$'' object IRAS~08572+3915 and Mrk~231 have $T_d$ of 74 K and 72 K, respectively, following the general trend expected of the dust temperature variation. On the other hand, two other dusty starbursts with characteristic $T_d\\ge70$ K, IRAS 15250+3609 and IRAS~05189$-$2524, have ``$z_{ph}$'' of +0.01 and +0.05, clearly demonstrating that there are other compensating effects and that the scatter in $T_d$ alone does not dictate the overall photometric redshift uncertainty. Regardless of the underlying causes, the ``photometric redshifts'' for 2/3 of all galaxies lie within $\\Delta z \\le 0.10$, and we estimate the collective uncertainty of the photometric redshift technique, including the variations in the radio and dust properties as well as the uncertainties in the $\\chi^2$ minimization process, is about 0.1 as long as this technique is applied to luminous dusty starburst galaxies only. Allowing for the ($1+z$) frequency folding of the Doppler effect, we estimate an overall uncertainty of $\\sigma_z \\sim 0.1 (1+z)$ with an upper bound in redshift uncertainty of about $0.3(1+z)$. In either case, this photometric redshift technique utilizing the radio-to-FIR dusty starburst SED represents a significant step forward, particularly at high redshifts, when compared with existing methods. The full potential of this method will be realized when the sources identified by several large deep, multi-frequency surveys planned in the immediate future (e.g. SIRTF Legacy Surveys) are analyzed together to reveal an accurate redshift distribution of luminous dusty galaxies at high redshift. \\bigskip" + }, + "0112/astro-ph0112238_arXiv.txt": { + "abstract": "We discovered that GSC 6554.559 is a previously unknown variable star, and named as TmzV429. We noticed that TmzV429 is identified with the IRAS-selected proto-planetary nebula (PPN), IRAS 08005-2356, which is undergoing a vigorous mass-loss episode. The analysis of photometric data suggests that TmzV429 a short-period pulsator, resembling a high-latitude yellow supergiant variable. TmzV429 is considered to be one of rare objects caught in the rapid course of PPN evolution, and shows one of the most striking mass-loss features among variable stars in the PPN stage. Since its evolutionary time-scale is estimated to be quite short ($\\sim$150 yrs), future observations of pulsations of TmzV429 is expected to provide an excellent opportunity to study the stellar evolution in real time. ", + "introduction": "Proto-planetary nebulae (PPNe) are objects in transition between the AGB stage and planetary nebula (PN) stage in stellar evolution (for a review, see \\cite{hri97}). PPNe are astrophysical objects not only important in studying the mass-loss from post-AGB stars and the formation of PNe, but also are considered to related to some of enigmatic high-latitude luminous yellow variables, such as RV Tau stars and UU Her stars (for a recent review, see \\cite{hrilu97}). TmzV429 (=GSC 6554.559)\\footnote{The permanent designation V510 Pup has been given.} is a variable star discovered by Takamizawa \\citep{tak99} The J2000.0 coordinates are 08\\h 02\\m 40\\s.71, $-$24\\deg 04\\arcm 42\\farcs 4. \\citet{tak99} reported small amplitude variations with a total photographic range of variability of 11.7--12.2. \\citet{tak99} originally suspected that this star is an semiregular variable. We discovered that this variable star, inconspicuous at the time of the variability announcement, is identified with a conspicuously mass-losing central star of a PPN, IRAS 08005-2356. We describe in this paper the analysis of our photometric data and the astrophysical implications of the present identification with a rapidly evolving PPN. ", + "conclusions": "We noticed that TmzV429 is identified with the PPN with a rapid mass-loss, IRAS 08005-2356 \\citep{sli91}. The object is also identified with an infrared source, MSX5C G242.3642+03.5822 \\citep{ega99}. The optical spectral classification by \\citet{sli91} is a late F-supergiant with prominent hydrogen emission lines. Together with Takamizawa's discovery of optical variability, the object seems to be classified as an high luminosity yellow supergiant variable (SRD-type in the General Catalogue of Variable Stars). \\citet{hek99} also reported a possible brightening by a several tenths of magnitude since 1986. This possible variation seems to be more likely attributed to shorter time-scale variations discovered by us. \\begin{figure} \\begin{center} \\includegraphics[angle=0,width=8cm]{lc.eps} \\caption{Light curve of TmzV429 drawn from the data in table \\ref{tab:obs}.} \\label{fig:lc} \\end{center} \\end{figure} \\begin{figure} \\begin{center} \\includegraphics[angle=0,width=7cm]{per1.eps} \\includegraphics[angle=0,width=7cm]{per2.eps} \\caption{(Left) Period analysis TmzV429. The most significant period of 28.9 d is marked with a tick. (Right) Folded light curve of TmzV429. The phase zero is taken arbitrarily.} \\label{fig:per} \\end{center} \\end{figure} We analyzed the original discovery data by \\citep{tak99} using the Phase Dispersion Minimization (PDM) method \\citep{ste78}. The result of period analysis are shown in figure \\ref{fig:per} (left panel). The period search was done for periods between 10 and 100 d. The range was limited mainly due to the data sampling, but covers most frequently met periods in low-mass, high luminosity, SRD-type variables. The strongest period between 10 and 100 d is 28.9 d. The period probably needs be treated with caution, because the period is close to the lunar month, and because of the possible intrinsic irregularity in such a variable. A rapid fading by 0.4 mag between JD 2450378 and 2450426, however, supports the existence of short-period variation with a period less than $\\sim$100 d. The folded light curve by this period is shown in figure \\ref{fig:per} (right panel). This result shows that the variability discovered by Takamizawa can be expressed by oscillations with a single, relatively short period. Although the possibility of a longer period can not be completely disregarded, the raw data (table \\ref{tab:obs}) suggest a short-period variation, rather than a period of hundreds of days to years. \\citet{sli91} reported some line features are similar to $\\rho$ Cas. The presently discovered variation, however, is not consistent with variations with a much longer period ($\\sim$300 d) as in massive $\\rho$ Cas-like variables. The star should be thereby regarded as a low-mass, post-AGB pulsator (e.g. \\cite{aik91}), which is consistent with the evolutionary stage \\citep{sli91} inferred from optical spectroscopy and IRAS observations. Although the number of observations is still limited, and the present analysis unavoidably suffers from a uncertainty, the present result suggests the existence of a low-amplitude ($\\sim$0.5 mag), relatively short-period pulsations in TmzV429, which are analogous to variations observed in some stellar components of other PPNe and in high galactic luminous yellow variables, such as RV Tau stars and UU Her stars. \\citet{sli91} suggested that the evolutionary time scale of this object is quite short ($\\sim$150 yrs). TmzV429 is thus one of rare objects caught in the rapid course of PPN evolution, and shows one of the most striking mass-loss features among variable stars in the PPN stage. Pulsations in such stars are a sensitive indicator of the evolution \\citep{aik91}, future observations of pulsations of this object will provide an excellent opportunity to study the {\\it stellar evolution in real time}." + }, + "0112/astro-ph0112148_arXiv.txt": { + "abstract": "{ The advent of robust, reliable and accurate higher order Godunov schemes for many of the systems of equations of interest in computational astrophysics has made it important to understand how to solve them in multi-scale fashion. This is so because the physics associated with astrophysical phenomena evolves in multi-scale fashion and we wish to arrive at a multi-scale simulational capability to represent the physics. Because astrophysical systems have magnetic fields, multi-scale magnetohydrodynamics (MHD) is of especial interest. In this paper we first discuss general issues in adaptive mesh refinement (AMR). We then focus on the important issues in carrying out divergence-free AMR-MHD and catalogue the progress we have made in that area. We show that AMR methods lend themselves to easy parallelization. We then discuss applications of the RIEMANN framework for AMR-MHD to problems in computational astophysics.} ", + "introduction": "Observations of various astrophysical systems such as proto-stars, novae, supernovae, the interstellar medium, galaxies and cosmology show that physical processes in these systems evolve in multi-scale fashion. To take but a single example from proto-star formation, once a Class 0 core forms out of the turbulent, magnetized, molecular gas in a molecular cloud, gravity causes the system to be largely decoupled from the rest of the turbulent flow. And yet, the decoupling is not complete. Large-scale magnetic fields cause the proto-star's angular momentum to be coupled to that of the external medium. Likewise, different parts of the dusty protostellar envelope are radiatively coupled to each other. The radiative coupling also sets the temperature of the gas, thereby determining its coupling with the magnetic field. As a result, we see the need for: (1) detailed representation of the micro-physical processes, (2) accurate representation of the system of equations that are of interest in computational astrophysics and (3) an ability to do (1) and (2) in a multi-scale fashion via adaptive mesh refinement (AMR). AMR is an elegant technique for concentrating mesh resolution in regions where an accurate answer is desired. Almost all astrophysical processes have magnetic fields and so an ability to represent MHD in multi-scale fashion is a {\\it sine qua non} for computational astrophysics. The present paper focusses on recent advances by the author that have made it possible to carry out AMR-MHD calculations. Ever since the first paper on AMR methods for fluid dynamics, see Berger and Colella (1989), it has been recognized that that higher order Godunov schemes play a central role in practical AMR simulations. The reasons are not difficult to see and the talks by Ryu (this conf.) on MHD, Ibanez and Koide (again this conf.) on relativistic flow bear testimony to the usefulness of these techniques in computational astrophysics. Such schemes were first formulated for hydrodynamics by vanLeer (1979), where their usefulness has been very well-documented by Woodward and Colella (1984). In recent years, they have been formulated for several other systems of interest in computational astrophysics wich include relativistic hydrodynamics, see Balsara (1994); radiation hydrodynamics, see Balsara (1999a,b,c); radiation MHD, see Balsara (1999d,e); relativistic MHD, see Balsara (2001a) and multidimensional radiative transfer, see Balsara (2001b). Perhaps the most vigorous evolution has taken place in MHD where Roe and Balsara (1996) designed the first complete MHD eigenvectors that were free of singularities; Brio and Wu (1988), Zachary, Malagoli and Colella (1994), Powell (1994), Dai and Woodward (1994), Ryu and Jones (1995) and Balsara (1998a) designed Riemann solvers for MHD; Dai and Woodward (1995), Ryu et al (1998) and Balsara (1998b) catalogued different forms of TVD schemes for MHD and Balsara and Spicer (1999), Dai and Woodward (1998), Ryu et al (1998), Londrillo and Del Zanna (2000) and Toth (2000) formulated divergence-free higher order Godunov schemes for MHD. The latter divergence-free formulations realize that the divergence of the magnetic field should remain exactly zero. Brackbill and Barnes (1980) and Brackbill (1985) have shown that violating the constraint leads to unphysical plasma transport orthogonal to the magnetic field as well as a loss of momentum and energy conservation. Powell et al (1999) did indeed formulate an AMR scheme for MHD that was not divergence-free only to find that the maximal errors ocurred on the finest meshes due to unphysical build-up of divergence. The finest meshes are the very meshes where one would have wanted the error to be minimal! Moreover, in accreting astrophysical flows, the divergence would flow with the fluid and build up in the very regions where one wants the most accurate answer! This prompted the present author to realize that only a scheme that was divergence-free on the entire AMR hierarchy would be adequate for astrophysical applications. These advances have been catalogued in detail in Balsara (2001c) and implemented in the RIEMANN framework for computational astrophysics. While Balsara (2001c) provides more mathematical details, the present paper focusses more on the intuitive ways of thinking about the subject. As a result, the two papers complement each other. In Section II we discuss algorithmic issues in AMR hydrodynamics. In Section III we discuss divergence-free AMR-MHD. In Section IV we show how AMR-MHD techniques have been parallelized. In Section V we discuss tests and applications. ", + "conclusions": "Several advances in divergence-free AMR-MHD and its application to astrophysics are reported here. We list them below: (1) A general strategy is presented for the time-update of the MHD system of equations on AMR hierarchies. (2) Just as Berger and Colella (1989) reduced the conservative time-update of the Euler equations on an AMR hierarchy to the application of a few simple steps, we have reduced the divergence-free time-update of the MHD equations on an AMR hierarchy to the application of a few simple steps. The steps have been summarized in Section III. (3) A significant advance has been made in the divergence-free reconstruction of vector fields. (4) Divergence-free prolongation of magnetic fields on an AMR hierarchy can be carried out via a very slight extension of the divergence-free reconstruction scheme mentioned in the previous point. (5) A divergence-free restriction strategy is presented. (6) An electric field correction strategy is presented which restores the consistency of electric fields at a fine-coarse interface in the AMR hierarchy. (7) Because of the above four points, the time-step can be sub-cycled on finer meshes without loss of the divergence-free property of the magnetic fields. (8) The above-mentioned innovations have been incorporated in the RIEMANN framework for parallel, self-adaptive computational astrophysics. Several stringent test problems have been presented and it is shown that the method presented in this paper for AMR-MHD is truly divergence-free. (9) Several very useful insights into astrophysical processes have been derived from the AMR-MHD simulations that have been presented." + }, + "0112/astro-ph0112462_arXiv.txt": { + "abstract": "We present spectroscopic observations of the X-ray transient XTE~J1118+480 acquired during different epochs following the 2000 March outburst. We find that the emission line profiles show variations in their double-peak structure on time scales longer than the 4.1 hr orbital period. We suggest that these changes are due to a tidally driven precessing disk. Doppler imaging of the more intense Balmer lines and the He{\\sc ii} $\\lambda$4686 line shows evidence of a persistent region of enhanced intensity superposed on the disk which is probably associated with the gas stream, the hotspot or both. We discuss the possible origins of the optical flux in the system and conclude that it may be due to a viscously heated disk. ", + "introduction": "X-ray Novae (XRNe) form a subclass of Low Mass X-ray binaries (LMXBs) which has provided us with some of the best stellar-mass black hole candidates (see the reviews by \\citealt{van95}, \\citealt{van98} and \\citealt{cha98}). These systems undergo occasional outbursts, with a recurrence time scale of decades, caused by an episode of intense mass transfer onto the compact primary via an accretion disk. A disk instability \\citep{can93,can98} is generally invoked as the triggering mechanism for the outburst. The XRN XTE~J1118+480 was discovered on 2000 March 29 by the Rossi X-Ray Timing Explorer All-Sky Monitor (RXTE ASM) at the beginning of a prolonged outburst \\citep{rem00}. Reanalysis of the ASM data showed a previous outburst episode during January 2-29 (MJDs~$51,545-51,572$; see Figure \\ref{fig1}). The optical counterpart was identified with a 12.9 magnitude star coincident with an 18.8 magnitude object in the USNO catalogues \\citep{uem00A}. \\citet{uem00B} claimed the detection of superhumps in the outburst light curve with a periodicity of 0.1709~d. The spectrum of the optical counterpart was typical of X-ray binaries in outburst \\citep{gar00}, but the X-ray flux was low for an X-ray transient \\citep{rem00}. A high inclination system \\citep{gar00} and/or a 'mini-outburst' state \\citep{hyn00} have been suggested to explain the anomalous X-ray flux. \\citet{dub01} presented the first constraints on the system parameters from the analysis of spectroscopic observations during the outburst. \\citet{mcc00,mcc01A} and \\citet{wag00,wag01} derived an orbital period that was shorter than the Uemura's determination (by $0.4-0.6~\\%$ respectively) from radial velocity measurements near quiescence. These authors determined the mass of the primary star to be $>6~M_\\odot$. This places XTE~J1118+480 among the dynamically confirmed black holes. In this paper we present the results from a systematic spectroscopic campaign. Data were collected on 21 nights during 2000 March-July; on four of these nights we obtained almost complete orbital phase coverage. ", + "conclusions": "Multi-epoch optical spectroscopic observations of XTE~J1118+480 yield valuable information about the properties and evolution of the accretion disk during outburst. The data reveal strong double-peaked H$\\alpha$, H$\\beta$, and He{\\sc ii} $\\lambda$4686 emission lines whose double-peak intensity varies on time scales longer than the 4.1 hr orbital cycle. These changes in the line profiles can be interpreted as resulting from a precessing eccentric disk around the compact primary. However, a search for periodic variability in the He{\\sc ii} $\\lambda$4686 emission line failed to find a modulation within the estimated precession period range. Therefore conclusive evidence for a precessing disk in the system is still required. Doppler tomograms display a bright bow-shaped emission in the -V$_X$, +V$_Y$ quadrant which we interpret as emission from the accretion gas stream, the hotspot or both. A more accurate ephemeris is needed to properly constrain the location of the emission within the binary system. A comparison with DNe in superoutburst shows that viscous dissipation in the disk may make a significant contribution to the optical luminosity during outburst." + }, + "0112/astro-ph0112181_arXiv.txt": { + "abstract": "We present a deuterium abundance analysis of the line of sight toward the white dwarf \\wdvd\\ observed with the Far Ultraviolet Spectroscopic Explorer (\\fuse). Numerous interstellar lines are detected on the continuum of the stellar spectrum. A~thorough analysis was performed through the simultaneous fit of interstellar absorption lines detected in the four \\fuse\\ channels of multiple observations with different slits. We excluded all saturated lines in order to reduce possible systematic errors on the column density~measurements. We report the determination of the average interstellar D/O and D/N ratios along this line of~sight at the 95\\% confidence level: D/O$\\;= 4.0 \\, (\\pm1.2) \\times 10^{-2}$; D/N$\\;= 4.4 \\, (\\pm1.3) \\times 10^{-1}$. % In conjunction with \\fuse\\ observations of other nearby sight lines, the results of this study will allow a deeper understanding of the present-day abundance of deuterium in the local interstellar medium and its evolution with time. ", + "introduction": "Deuterium is believed to be produced in appreciable quantities only in primordial Big Bang nucleosynthesis (BBN) and destroyed in stellar interiors (\\eg, Epstein et al.~\\citealp{epstein76}); it is thus a key element in cosmology and in galactic chemical evolution (see \\eg~Vangioni-Flam \\& Cass\\'e~\\citealp{flam95}; Prantzos~\\citealp{prantzos96}; Scully et al.~\\citealp{scully97}). The primordial abundance of deuterium is one of the best probes of $\\Omega_Bh^2$, the baryonic density of the Universe divided by the critical density. The abundance of deuterium relative to hydrogen (D/H) is expected to decline during Galactic evolution at a rate that is a function of the star formation rate; standard models predict a factor of 2 to 3 decrease in the deuterium abundance in 15 Gyrs (see \\eg, Galli et al.~\\citealp{galli95}; Tosi et al.~\\citealp{tosi98}). Hence, any abundance of deuterium measured at any metallicity should provide a lower limit to the primordial deuterium abundance. This picture is essentially constrained by deuterium abundance measurements at look-back times of $\\sim14$~Gyrs (primordial intergalactic clouds), 4.5~Gyrs (protosolar), and 0.0~Gyrs (\\ism). Although the evolution of the deuterium abundance seems to be qualitatively understood, measurements of D/H at similar redshifts show some dispersion and indicate that additional processes may be important in controlling the abundance of deuterium. That fact has led to the development of non-standard models, which propose, for example, larger astration factors (\\eg, Vangioni-Flam et al.~\\citealp{flam94}) or non-primordial deuterium production (see~\\eg, Lemoine et al.~(\\citealp{lemoine99}) for a~review). Up to now, the \\ism\\ is the astrophysical site that has allowed the most comprehensive investigations of deuterium abundances. Deuterium has been observed in the \\ism\\ using different methods: radio measurements of its 92~cm hyperfine transition (\\eg, Blitz \\& Heiles~\\citealp{blitz87}; Chengalur et al.~\\citealp{chenga97}), observations of deuterated molecules (\\eg, Lubowich et al.~\\citealp{lubo00}; Ferlet et al.~\\citealp{ferlet00}), Balmer series analyses (H\\'ebrard et al.~\\citealp{hebrard00}), and Lyman series absorption (see Moos et al.~\\citealp{moos01}). Of these, the most accurate measurements have been obtained through Lyman absorption-line observations in the far-ultraviolet (far-UV) spectral range. By observing hydrogen and deuterium directly in their atomic form, far-UV Lyman series absorption-line measurements provide accurate column density determinations that are not dependent on ionization or chemical fractionation effects. The first measurement of the abundance ratio \\dshism\\ was reported by Rogerson \\& York~(\\citealp{ry73}) for the line of sight to $\\beta$~Cen, using {\\it Copernicus}: \\dshism$=1.4\\pm0.2\\times10^{-5}$. Since then, many other \\dshism\\ measurements have been performed using different instruments ({\\it Copernicus}, {\\it IUE}, {\\it IMAPS}, {\\it HST}) for other sight lines, and the values obtained show significant dispersion around the above~value. For example, an average value \\dshism$=1.50(\\pm0.10)\\times10^{-5}$ ($1\\,\\sigma$) has been derived for the Local Interstellar Cloud (Lallement \\& Bertin~\\citealp{lallement92}) by Linsky~(\\citealp{linsky98}) from the comparison of 12 nearby sight lines, but studies of several lines of sight revealed values outside this range (\\eg,~Laurent et al.~\\citealp{laurent79}; York~\\citealp{york83}; Allen et al.~\\citealp{allen92}; Vidal-Madjar et al.~\\citealp{avm98}; H\\'ebrard et al.~\\citealp{hebrard99}; Jenkins et al.~\\citealp{jenkins99}; Sonneborn et al.~\\citealp{sonneborn00}). This dispersion may result from spatial variations due to some unknown physical processes or underestimation of systematic errors. There is still considerable debate over these two interpretations, and the final resolution of the issue may have implications for understanding the physics of the interstellar medium, as well as the baryonic density inferred from D/H measurements. Measurements of D/H in the intergalactic medium are also sparse and do not agree with each other (\\eg,~Webb et al.~\\citealp{webb97}; Burles \\& Tytler~\\citealp{burles98a}; \\citealp{burles98b}), so an accurate determination of the primordial deuterium abundance has also proven to be a difficult quantity to measure. Moreover, recent studies of the anisotropy of the cosmic microwave background (CMB), which permits evaluations of the baryonic density independent of those obtained through deuterium measurements (see, \\eg, de~Bernardis et al.~\\citealp{bernardis00}; Jaffe et al.~\\citealp{jaffe01}), imply higher values of $\\Omega_Bh^2$ than those implied by the abundance studies, corresponding to a primordial D/H value of the order of the \\dshism~values (see however de~Bernardis et al.~\\citealp{bernardis01}; Stompor et al.~\\citealp{stompor01}; Pryke et al.~\\citealp{pryke01}). An accurate determination of the interstellar deuterium abundance is one of the main objectives of the Far Ultraviolet Spectroscopic Explorer (\\fuse), which was successfully launched on 1999 June 24 (Moos et al.~\\citealp{moos00}). In this paper, we present new measurements of deuterium abundances obtained with \\fuse\\ toward the white dwarf \\wdvd. This paper is part of a series dedicated to \\fuse\\ measurements of deuterium interstellar abundances toward BD\\,+28$^{\\circ}$\\,4211 (Sonneborn et al.~\\citealp{sonneborn01}), WD$\\,$1634$-$573 (Wood et al.~\\citealp{wood01}), WD$\\,$0621$-$376 (Lehner et al.~\\citealp{lehner01}), G191$-$B2B (Lemoine et al.~\\citealp{lemoine01}), HZ$\\,$43A (Kruk et al.~\\citealp{kruk01}), and Feige~110 (Friedman et al.~\\citealp{friedman01}). A major goal of this program is to determine to what extent and on what scales variations in the D/H ratio occur. Moos et al.~(\\citealp{moos01}) present an overview of these first results. White dwarfs are ideal targets for the measurement of the interstellar deuterium abundance (Lemoine et al.~\\citealp{lemoine96}); they may be chosen close to the Sun (so that the column densities are not too high and the velocity structure of the line of sight is not too complex), and they exhibit relatively smooth UV continua (which allow detections of lines from many species). Prior to this series of papers, measurements have been published for \\dshism\\ toward three white dwarfs: G191$-$B2B (Lemoine et al.~\\citealp{lemoine96}; Vidal-Madjar et al.~\\citealp{avm98}; Sahu et al.~\\citealp{sahu99}), Sirius-B (H\\'ebrard et al.~\\citealp{hebrard99}), and Feige~24 (Vennes et al.~\\citealp{vennes00}). The \\fuse\\ observations of \\wdvd\\ and the data processing are presented in Sect.~\\ref{Observations_and_data_processing}, and the details of the analysis in Sect.~\\ref{Data_Analysis}. The results are reported in Sect.~\\ref{Results} and discussed in Sect.~\\ref{Discussion}. ", + "conclusions": "\\label{Conclusion} We have presented \\fuse\\ observations of the 905-1187~\\AA\\ spectrum of the white dwarf \\wdvd. The column densities of several ions were measured through simultaneous fits to the numerous unsaturated absorption lines detected in the four channels and through the three apertures. In particular, the ratios D/O$\\;= 4.0 \\, (\\pm1.2) \\times 10^{-2}$ and D/N$\\;= 4.4 \\, (\\pm1.3) \\times 10^{-1}$ (\\ds\\ error bars) were measured. This result is discussed by Moos et al.~(\\citealp{moos01}) together with other early \\fuse\\ deuterium results. Observations of the \\lya\\ absorption toward \\wdvd\\ would help to constrain the D/H ratio along this sight line. Upcoming \\fuse\\ observations will help to determine if the D/H and D/O ratios vary locally. The answer to this question is critical to improving our understanding of the abundance of deuterium." + }, + "0112/astro-ph0112294_arXiv.txt": { + "abstract": "We report the results of our radial-velocity monitoring of spectroscopic binary systems in a sample of X-ray sources from the ROSAT All Sky Survey south of the Taurus-Auriga star-forming region. The original sample of $\\sim$120 sources by Neuh\\\"auser et al.\\ was selected on the basis of their X-ray properties and the visual magnitude of the nearest optical counterpart, in such a way as to promote the inclusion of young objects. Roughly 20\\% of those sources have previously been confirmed to be very young. We focus here on the subset of the original sample that shows variable radial velocities (43 objects), a few of which have also been flagged previously as being young. New spectroscopic orbits are presented for 42 of those systems. Two of the binaries, RXJ0528.9$+$1046 and RXJ0529.3$+$1210, are indeed weak-lined T~Tauri stars likely to be associated with the $\\lambda$~Ori region. Most of the other binaries are active objects of the RS~CVn-type, including several W~UMa and Algol systems. We detect a strong excess of short-period binaries compared to the field, and an unusually large fraction of double-lined systems. This, along with the overall high frequency of binaries out of the original sample of $\\sim$120 sources, can be understood as a selection effect since all these properties tend to favor the inclusion of the objects in a flux-limited X-ray survey such as this by making them brighter in X-rays. A short description of the physical properties of each binary is provided, and a comparison with evolutionary tracks is made using the stellar density as a distance-independent measure of evolution. We rely for this on our new determinations of the effective temperature and projected rotational velocities of all visible components of the binaries. A number of the systems merit follow-up observations, including at least 4 confirmed or probable eclipsing binaries. One of these, RXJ0239.1$-$1028, consists of a pair of detached K dwarfs and may provide for a potentially important test of stellar evolution models once the absolute dimensions of the components are determined. ", + "introduction": "\\label{secintro} With the launch of the ROSAT satellite and of its predecessor, the Einstein Observatory, a new dimension was added to the study of star-forming regions that has enabled us to probe the higher energies released by young stellar objects. Many X-ray sources associated with known classical T~Tauri stars (cTTS) or weak-line T~Tauri stars (wTTS) were detected by these space missions in molecular cloud regions that have traditionally been thought of as being the natural habitat of low-mass pre-main sequence (PMS) stars. Additional wTTS also in the immediate vicinity of the molecular gas were discovered, nearly doubling the known population of young stars of this kind. But large numbers of similar objects have also turned up \\emph{around} these regions, far removed from the presently known cloud material. Systematic studies to search for young objects over large areas of the sky based on the (spatially unbiased) ROSAT All Sky Survey (RASS) have been carried out by \\markcite{a95}Alcal\\'a et al.\\ (1995) in the Chamaeleon region, \\markcite{w96}Wichmann et al.\\ (1996) in the central region of the Taurus-Auriga association, \\markcite{a96}Alcal\\'a et al.\\ (1996) in Orion, \\markcite{k97}Krautter et al.\\ (1997) in Lupus, \\markcite{n95a}Neuh\\\"auser et al.\\ (1995a; hereafter N95a) and \\markcite{m97}Magazz\\`u et al.\\ (1997; hereafter M97) south of the Taurus-Auriga region, and \\markcite{n00}Neuh\\\"auser et al.\\ (2000) in Corona Australis. These studies, mostly carried out at relatively low spectral resolution, have identified several hundred candidates that also have all the appearance of wTTS and their sheer numbers could have a significant impact on our understanding of issues such as the efficiency of the process of star formation. Follow-up studies with high spectral resolution to confirm the characteristic spectroscopic signatures of T~Tauri stars among these candidates (late spectral type along with H$\\alpha$ emission and strong \\ion{Li}{1}~$\\lambda$6708 absorption) have also been performed, and a fraction of the objects ($\\sim$20-90\\%, depending on the region; e.g., \\markcite{c97}Covino et al.\\ 1997, \\markcite{w99}Wichmann et al.\\ 1999, \\markcite{w00}2000) indeed seem to be very young (1-10~Myr). Kinematic studies to investigate the association of these objects with their parent clouds have provided important complementary information (\\markcite{n95a}N95a; \\markcite{n97}Neuh\\\"auser et al.\\ 1997, hereafter N97; \\markcite{a00}Alcal\\'a et al.\\ 2000; \\markcite{w00}Wichmann et al.\\ 2000). Radial velocities for more than 100 X-ray sources in the region south of the Taurus-Auriga molecular clouds were reported by \\markcite{n95a}N95a and \\markcite{n97}N97. In the course of that high-resolution spectroscopic survey we discovered several dozen spectroscopic binaries, three of which have been identified as pre-main sequence stars in earlier studies. Such systems are extremely interesting given that orbits have been determined for only about 40 PMS systems (\\markcite{m94}Mathieu 1994; \\markcite{m01}Melo et al.\\ 2001), and the study of their properties can provide valuable insights into the process of star formation. In special cases these binaries allow one to determine the absolute masses of the components ---the most basic property of a star--- and can serve to test models of stellar evolution for young stars that are virtually unconstrained by the observations so far. In this paper we present the results for the binary population of the sample of X-ray sources observed by \\markcite{n95a}N95a and \\markcite{n97}N97, including new orbital solutions for 42 objects and a detailed discussion of each system. We also report determinations of the effective temperature and projected rotational velocity for all visible components, which aid in establishing the nature of these objects by comparison with recent stellar evolution models. As expected from their detection in X-rays by ROSAT, all of them are fairly active and this is reflected in some of their overall properties such as the orbital period distribution, which we compare with that for normal solar-type stars. Perhaps not surprisingly, the great majority of the systems are most likely of the RS~CVn type, and a few belong to the Algol or W~UMa class of interacting binaries. For at least two of our stars the spectroscopic and dynamical evidence suggests that they are bona-fide PMS objects, one of them possibly being as young as 1-2~Myr. These are potentially important systems that merit further study. ", + "conclusions": "\\subsection{Lithium strengths and PMS status} \\label{secpms} The original selection of the \\markcite{n97}N97 sample, while designed to favor the detection of young stars in a flux-limited X-ray survey (RASS), has obviously not been completely successful as attested by the large fraction of (post-main sequence) RS~CVn systems among our binaries. Nevertheless, a handful of our program stars do have detectable \\ion{Li}{1}~$\\lambda$6708 absorption, one of the key diagnostics of youth used to support PMS status. Because the strength of the Lithium line depends quite strongly on the effective temperature, and even on the rotational velocity of the star (see, e.g., \\markcite{s93}Soderblom et al.\\ 1993; \\markcite{mc96}Mart{\\'\\i}n \\& Claret 1996), both of these parameters must be accounted for to evaluate the significance of the Li signature. Even with this, additional information on the systems can also be very helpful, such as that provided by our orbital solutions in \\S\\ref{secresults} and \\S\\ref{seccomments}. In Figure~\\ref{figli} we represent the Li measurements as a function of effective temperature for the 10 objects in our sample with equivalent widths larger than 0.1~\\AA, as listed in Table~\\ref{tabsample}. For stars with more than one measurement we have adopted the average. Six systems are double-lined, but only two have separate measurements for the two components. For the others, we have adopted the temperature of the brighter component and assumed that the Li measurement corresponds largely to that star. The measurements for all double-lined systems have been corrected for the dilution factor due to binarity, using our temperature determinations and approximate light ratios at 6708~\\AA\\ estimated from our measurements at 5187~\\AA. \\placefigure{figli} The segmented lines in the Figure~\\ref{figli} represent the upper envelope of the Li distributions for two young clusters often used for comparison purposes --- IC~2602 (age $\\sim 35$~Myr), and the Pleiades (age $\\sim 100$~Myr) --- following \\markcite{n97}N97. For the latter cluster, which has been studied in considerably more detail, the rapid rotators ($v \\sin i > 15\\kms$) display somewhat lower Li depletion (larger equivalent widths) than the slow rotators, as represented by the solid and long-dashed lines, respectively. Three of our systems deserve special comment. The cooler component of [24]~RXJ0422.9$+$0141 lies slightly above the upper envelope for IC~2602, suggesting the possibility of a PMS status. However, the dynamical and physical information discussed in \\S\\ref{seccomments} seems to rule this out (see also below), since the larger mass for this cooler component indicates that this star must have evolved to become a subgiant or a giant. We conclude therefore that this is a post-main sequence system of the RS~CVn type. Studies by \\markcite{prg92}Pallavicini, Randich \\& Giampapa (1992), \\markcite{ff93}Fern\\'andez-Figueroa et al.\\ (1993), \\markcite{r94}Randich, Giampapa \\& Pallavicini (1994) and others have shown that active systems of the RS~CVn type and other categories of active objects occasionally show unexpectedly high levels of Li absorption compared to inactive stars of similar spectral types. While the mechanism responsible for the excess Li abundance is not yet clear, the case of [24]~RXJ0422.9$+$0141 appears consistent with those findings. The two components of [41]~RXJ0528.9$+$1046 show significant excess Li absorption, and the supplementary evidence discussed in \\S\\ref{seccomments} strongly supports its classification as a weak-lined T~Tauri system. The Li strength measured for [42]~RXJ0529.3$+$1210 puts it essentially at the upper envelope of IC~2602, but this may be only a lower limit since there are hints in our spectra of light from one or perhaps two additional components. This extra light would cause the true equivalent width of the primary to be underestimated, though it is difficult at the moment to determine by how much. Based on this and the strong suggestion of association with the $\\lambda$~Ori region discussed in \\S\\ref{seccomments}, we consider it quite likely that this system is also very young. Thus, two of our program stars display properties indicating that they are PMS objects with ages most likely under $\\sim$30~Myr. The fact that they are binaries makes them especially interesting, given that only about three dozen such objects have had their spectroscopic orbital elements determined. Furthermore, [42]~RXJ0529.3$+$1210 is among only a handful with periods long enough that the components may be spatially resolved in the near future using large ground-based interferometers observing in the infrared. At these wavelengths the contrast between the primary and secondary components should be much more favorable than in the optical. At an assumed distance of 450~pc for [42]~RXJ0529.3$+$1210 (\\markcite{dm01}Dolan \\& Mathieu 2001) the angular semimajor axis of the pair is estimated to be about 3~mas. But because of the large eccentricity of its orbit, the maximum separation can be as large as twice this value. \\subsection{A comparison with evolutionary tracks} \\label{secdensity} In \\S\\ref{seccomments} we made use of the measured $v \\sin i$ values for the short-period double-lined binaries in our sample that have circular orbits, in order to provide estimates of the radii and surface gravities of the components and help in understanding their nature. In a sense the rotational velocities have thus been used as a sensitive measure of evolution, under the reasonable hypothesis that tidal forces have already synchronized the rotation of the stars and aligned their axes with that of the orbit. Those estimates are only lower limits, though, because of the unknown inclination angles. In this section we take this approach one step further and eliminate the dependence on $i$ to provide a physical magnitude that may be compared directly with predictions from theory for normal stars. Eq.(1) and eq.(2) give the quantities $M \\sin^3 i$ and $R \\sin i$ directly in terms of observable properties ($P$, $K$, and $v \\sin i$). The minimum masses and minimum radii have different dependences on $\\sin i$, but the ratio $M \\sin^3 i/(R \\sin i)^3$ removes the dependency, and happens to represent the \\emph{density} of a star ($M/R^3$) in terms of the solar density. Expressed directly as a function of the measurable parameters, we have \\begin{equation} \\log {\\rho_{A,B}\\over\\rho_{\\sun}} = -1.8718 - 2\\log P + 2\\log(K_A+K_B) + \\log K_{B,A} - 3\\log V^{\\rm rot}_{A,B}~, \\end{equation} where all quantities are given in their usual units. The density is intimately related to the internal structure and evolutionary status of a star, and the possibility of using this ``observable\" property in suitable double-lined systems to compare directly with stellar evolution models is often overlooked\\footnote{For a recent application of this idea, see \\markcite{q00}Quast et al.\\ (2000).}. The density changes by more than two orders of magnitude over the life of a normal star, so that at the very least this allows us (in conjunction with the measured effective temperatures) to tell whether the components are still on the main sequence or whether they are evolved. As it turns out, the discriminating power of the density is quite significant although it depends greatly on the precision of the rotational velocities, as seen from the factor of 3 in the last term of eq.(6). In Figure~\\ref{figdens} we show the components of all the double-lined systems in our sample that have sufficiently well determined rotational velocities, in a diagram of $\\log \\rho/\\rho_{\\sun}$ vs.\\ $\\log T_{\\rm eff}$. In all cases the periods are short or there is otherwise good reason to believe that the assumption of synchronization and alignment is valid, as described in the individual notes for each system. In each binary a dotted line connects the primary components (dark symbols) with the secondaries (open symbols, with the identification number). The observations are compared with theoretical isochrones for the main-sequence and post-main-sequence phase based on the models by \\markcite{y01}Yi et al.\\ (2001), for solar metallicity and ages of 1~Gyr, 3~Gyr, 10~Gyr and 15~Gyr. \\placefigure{figdens} Almost all of the secondary components are seen to lie near the main sequence, while roughly half of the primaries are evolved and lie at the base of the giant branch or above. A gap is seen in the diagram between the evolved and un-evolved stars that coincides with the locus of objects in the rapid subgiant phase, where fewer stars are expected at any given time. For the most part the two components of each binary are consistent with being on a single isochrone, given the errors, the most obvious exceptions being [12]~RXJ0309.1$+$0324N and [41]~RXJ0528.9$+$1046, which seem to be aligned in a direction roughly perpendicular to the isochrones. The first of these is most likely a W~UMa-type system, and its components should therefore not be expected to behave as single stars since they are in close mechanical and thermal contact. The evidence presented earlier for the other system indicates that it is likely to be in the PMS stage (see below). But for all the other objects, we note that the location of each star is fully consistent with our interpretation in \\S\\ref{seccomments}. This is expected since the same hypothesis was used, but the evolutionary status of each binary is more straightforward to see and to compare with the models in this figure. The densities of the primary components in all the binaries are typically smaller than those of the secondaries, as dictated by stellar evolution, except in the case of [30]~RXJ0441.9$+$0537, which is the highest object on the diagram in Figure~\\ref{figdens}. In this system (which is the ``cool Algol\" referred to earlier) mass transfer has caused a reversal of the mass ratio, so that the currently more massive star is actually the one that was originally the secondary (see \\markcite{t98}Torres, Neuh\\\"auser \\& Wichmann 1998). \\placefigure{figdenspms} As mentioned earlier, a few of our systems exhibit moderately strong \\ion{Li}{1}~$\\lambda$6708 absorption that may indicate an early stage of evolution (PMS), as opposed to a post-main sequence status. Six of them are double lined. It is quite illuminating to represent these systems in a $\\log\\rho/\\rho_{\\sun}$ vs.\\ $\\log T_{\\rm eff}$ diagram such as Figure~\\ref{figdens}, but to compare them instead with evolutionary tracks appropriate for the PMS stage. This is done in Figure~\\ref{figdenspms}, where the isochrones and evolutionary tracks are from the models by \\markcite{s97}Siess, Forestini \\& Dougados (1997) (see also \\markcite{s00}Siess, Dufour \\& Forestini 2000) for solar metallicity. The components of [5]~RXJ0219.7$-$1026, [14]~RXJ0339.6$+$0624, [16]~RXJ0343.6$+$1039, and [24]~RXJ0422.9$+$0141 are clearly inconsistent with being on the same isochrone, and this agrees with our assessment in \\S\\ref{seccomments} that they are not likely to be in the PMS phase based on dynamical and physical evidence (e.g., the mass ratio and luminosity ratio), or based on their relatively weak Li line in most cases (\\S\\ref{secpms}). Objects [35]~RXJ0444.4$+$0725 and [41]~RXJ0528.9$+$1046, on the other hand, do seem to conform to the isochrones in this diagram, within the errors. The age inferred for the first of these is not particularly young, and its properties (including the weak Li absorption) are consistent with it being a ZAMS star. The location in the $\\log\\rho/\\rho_{\\sun}$ vs.\\ $\\log T_{\\rm eff}$ plane is also consistent with the main-sequence isochrones shown in Figure~\\ref{figdens}, so that the case for a PMS status is not very strong in this system. More interesting is the agreement of [41]~RXJ0528.9$+$1046 with the PMS isochrones in Figure~\\ref{figdenspms}, while running \\emph{across} those in Figure~\\ref{figdens}. This, along with the Li strengths and kinematic information reported earlier, makes a more compelling case for this system being a much younger object. The absence of H$\\alpha$ emission reported by \\markcite{m97}M97 and by \\markcite{dm01}Dolan \\& Mathieu (2001) place it in the wTTS category. The age we derive for this SB2 from Figure~\\ref{figdenspms} is approximately 2~Myr, consistent with the estimate of 2.3~Myr by \\markcite{dm01}Dolan \\& Mathieu (2001) based on a different set of models (\\markcite{ps99}Palla \\& Stahler 1999). Use of yet another set of PMS isochrones by \\markcite{y01}Yi et al.\\ (2001) yields a younger age of about 1~Myr. The absolute masses for the components from this figure ($M_A = 1.96$~M$_{\\sun}$ and $M_A = 1.46$~M$_{\\sun}$) are not quite consistent with the spectroscopic mass ratio, which is measured much more accurately. This is most likely due to errors in the temperature determinations. In addition, the absolute masses are strongly model-dependent. For example, the values inferred from the PMS evolutionary tracks by \\markcite{y01}Yi et al.\\ (2001) are systematically lower by 0.4-0.5~M$_{\\sun}$, which is rather significant. Nevertheless, one may derive an average inclination angle for the orbit that is roughly 33-37\\arcdeg, which rules out eclipses, as indicated in \\S\\ref{seccomments}. It is worth pointing out, to conclude this section, that a diagram such as Figure~\\ref{figdenspms} can be a very useful tool in the field of young stars, specifically for the case of double-lined spectroscopic binaries with short periods and circular orbits. It is similar to the more widely used H-R diagram, but has the important advantage that it is completely independent of any assumption on the distance, a common criticism that is painfully familiar to all users of the $\\log L/L_{\\sun}$ vs.\\ $\\log T_{\\rm eff}$ diagram. \\subsection{Period distribution and binary frequency} \\label{secpdistrib} In the course of solving for the orbits of the binaries in the \\markcite{n97}N97 sample we had the distinct impression that the fraction of short-period systems was unusually high. Due to the nature of the target list, which was selected to include objects that are strong in X-rays, a bias towards short periods is actually expected from the fact that those systems will tend to be rotating faster because they are tidally locked, and thus will have enhanced activity and should be detected more easily by ROSAT. As it turns out, more than half of our binaries have orbital periods of less than 5 days. In addition, we noticed two other trends that were not expected for field stars, but in retrospect can also be understood in terms of selection effects. One is that the frequency of binaries among the ROSAT sources south of Taurus seemed higher than normal (43 spectroscopic binaries with orbits, out of a sample of 121 stars), and the other is that the fraction of double-lined systems (as opposed to single-lined binaries) also seemed high: 26 SB2 systems out of 43 binaries, or 60\\%. The detection of spectroscopic binaries in our sample is by no means complete as a function of orbital period. The majority of the stars were observed over an interval of 3 to 5 years, and we estimate conservatively that we have detected most of the systems with orbital periods up to about 1~yr, although we also found 3 binaries with even longer periods. The incompleteness due to small velocity amplitudes (low-mass companions, small inclinations, or both) is difficult to quantify, but it is unlikely to be very large for this sample and will not be considered here. As a benchmark for studying the frequency and period distribution of the binaries in our X-ray sample we have used the results for the solar-type stars in the solar neighborhood by \\markcite{dm91}Duquennoy \\& Mayor (1991). In Figure~\\ref{figpdist} we show the period distribution for the 43 binaries in our sample compared to the corresponding distribution for the field, on a logarithmic scale. The normalization was computed by requiring that the integral of the smooth curve up to periods of 1~yr (our completeness level, represented with a dashed line) is equal to the number of binaries detected in our sample up to that period (40). The hatched area of the histogram represents the double-lined systems, and the open area added above is for the SB1 binaries. \\placefigure{figpdist} The observed distribution of X-ray binaries is sharply peaked at short periods (median $\\log P = 0.67$, or a period of 4.7~days) and shows an excess over the field distribution, as anticipated above. The frequency of binaries in our sample (40 systems detected up to our completeness limit of 1~yr, out of 121 objects studied) is 33\\%~$\\pm$~5\\%. This is nearly double the fraction of binaries in the field up to this period, which is 17\\%~$\\pm$~3\\%. If we choose a more conservative completeness limit for our sample of, say, 100~days, the excess of binaries would be even greater since the number of systems detected in our sample up to 100~days remains the same, yet the binary frequency in the field decreases to 11\\%~$\\pm$~2\\%. The increased fraction of binaries compared to the field is understood as a selection effect, since binaries will tend to be brighter than single stars in X-rays, not only because they are composed of two objects, but also because of the increased activity if the period is short. They will thus tend to be promoted into the (flux-limited) sample more easily compared to single stars. A similar effect can explain why SB2 systems dominate over SB1 systems. \\subsection{Activity} \\label{secactivity} Correlations between some of the physical properties of the stars and a number of activity indicators might in principle be expected to exist in a sample such as ours, composed largely of RS~CVn-type or similarly active systems. For example, we searched for a dependence of the X-ray hardness ratios HR1 and HR2 and also the X-ray flux of our binaries (using data as reported by \\markcite{m97}M97) with parameters such as the orbital period, the effective temperature, the projected rotational velocity, and the surface gravity. No significant correlations were found, possibly because essentially all of our objects are active by virtue of their selection. For the double-lined systems where there is ambiguity as to which component's parameters to use, we considered alternatively those of the coolest component, the more massive component, and also the visually brightest component, with null results in all cases. No differences were found between the X-ray properties of the single-lined binaries and the double-lined binaries, with the possible exception of a hint that the SB2 systems may have systematically larger HR2 indices (by about 0.10-0.20), indicative of slightly harder X-ray emission. Similarly, examination of correlations between the H$\\alpha$ emission as given by \\markcite{m97}M97 (equivalent width) and the orbital or stellar characteristics $\\log P$, $v \\sin i$, or $\\log g$ indicated no dependence. There is, however, a clear dependence with effective temperature, in the sense that cooler systems tend to have stronger H$\\alpha$ emission. This is consistent with the fact that stars with lower temperatures are typically more chromospherically active. The correlation is shown in Figure~\\ref{fighalpha}a for the single-lined systems in our sample (triangles). Negative equivalent widths indicate emission, while positive values represent absorption (which may be partially filled-in in some cases). Although there appears to be a similar correlation for the double-lined systems (not shown), the scatter is much larger because of the confusion as to which component the emission measured by \\markcite{m97}M97 corresponds to. \\placefigure{fighalpha} It is of interest to compare the trend for the single-lined binaries with that for the single stars in the \\markcite{n97}N97 sample, which we have represented by crosses in Figure~\\ref{fighalpha}a. The temperatures for these single stars were derived using the same technique we have applied in this paper. The increase in H$\\alpha$ equivalent width for the cooler objects is again obvious. The slope of the correlation is essentially the same in both cases, but the binaries appear to lie slightly higher on the diagram. To test the statistical significance of this we carried out a linear fit to all the data (SB1 $+$ singles) and examined the residuals from this fit separately for the SB1 systems and the single stars. The slope of this correlation is $d({\\rm H}\\alpha)/dT_{\\rm eff} = +0.236 \\pm 0.026$~\\AA\\ per 100~K. Figure~\\ref{fighalpha}b shows the distributions of residuals. The SB1 systems have equivalent widths that are $0.56 \\pm 0.18$~\\AA\\ more negative, on average, indicating stronger emission (a 3.1$\\sigma$ effect). To examine whether this excess emission compared to the single stars might be due to a difference in rotational velocities, we show in Figure~\\ref{figvsini}a the $v \\sin i$ distributions for the single stars and for the SB1s. Despite expectations (tidal locking in close binaries), the distributions are not statistically different, as indicated by the Kolmogorov-Smirnov test. As it turns out, a significant fraction of the binaries have long orbital periods, and are therefore not yet affected by tidal effects. The reason for the slightly stronger H$\\alpha$ emission for the binaries in our sample is thus not obvious. \\placefigure{figvsini} The vertical scatter in Figure~\\ref{fighalpha}a, particularly for the single stars, does appear to be related to the rotation of the stars. In Figure~\\ref{figvsini}b we show the residuals from the H$\\alpha$ vs.\\ $T_{\\rm eff}$ correlation as a function of the rotational velocity of the objects, separately for the single stars and the binaries. The residuals for both types of objects have been corrected for the systematic difference found above. The trend for the single stars goes in the direction expected, namely, the fast rotators show extra emission and the slow rotators show less. There appears to be no such trend for the single-lined binaries, although the sample is small." + }, + "0112/astro-ph0112541_arXiv.txt": { + "abstract": "A magnetic dynamo experiment is under construction at the New Mexico Institute of Mining and Technology. The experiment is designed to demonstrate in the laboratory the $\\alpha\\omega$ magnetic dynamo, which is believed to operate in many rotating and conducting astrophysical objects. The experiment uses the Couette flow of liquid sodium between two cylinders rotating with different angular velocities to model the $\\omega$--effect. The $\\alpha$--effect is created by the rising and expanding jets of liquid sodium driven through a pair of orifices in the end plates of the cylindrical vessel, presumably simulating plumes driven by buoyancy in astrophysical objects. The water analog of the dynamo device has been constructed and the flow necessary for the dynamo has been demonstrated. Results of the numerical simulations of the kinematic dynamo are presented. The toroidal field produced by the $\\omega$--effect is predicted to be $B_{\\phi} \\simeq (R_m/2\\pi) B_{poloidal}\\simeq 20 \\times B_{poloidal}$ for the expected magnetic Reynolds number of $R_m \\sim 120$. The critical rate of jets necessary for the dynamo self-excitation is predicted from the calculations to be a pair of plumes every 4 revolutions of the outer cylinder. For reasonable technical limitations on the strength of materials and the power of the drive, the self-excitation of the dynamo appears to be feasible. ", + "introduction": "\\label{sec_dyn} Fig.~\\ref{flux_schem}(a) shows how the $\\alpha \\omega$ dynamo works in the New Mexico Dynamo Experiment and Fig.~\\ref{flux_schem}(b) shows how the $\\alpha\\omega$ dynamo works in the accretion disk forming the massive, central galactic black hole. In Fig.~\\ref{flux_schem}(a) differential rotation is established in the liquid sodium between two rotating cylinders as limiting stable Couette flow, $\\Omega \\propto 1/R^2$, by driving $\\Omega_{1} = 4 \\Omega_{0}$ where $R_{0} = 2 R_{1}$ and for the disk Fig.~\\ref{flux_schem}(b) as Keplerian rotation, $\\Omega \\propto 1/ R^{-3/2}$, around the central mass or the black hole. This differential rotation wraps up the radial component of an initial poloidal field Fig.~\\ref{flux_schem}(b(A)) either made with coils or an infinitesimally small, $< 10^{-19} $ G seed field from density structure at decoupling. The resulting toroidal field becomes stronger than the initial poloidal field Fig.~\\ref{flux_schem}(b(B)) by $B_{toroidal} / B_{poloidal} = n_{\\Omega} B_{poloidal}$, where in equilibrium with resistive decay, the limiting number of turns becomes $n_{\\Omega} \\simeq R_{m,\\Omega}/2\\pi$, and $R_{m,\\Omega} = v_0 (R_0 - R_1)/\\eta$ is the magnetic Reynolds number. This multiplication factor depends upon the resistivity of liquid sodium or for the disk, upon the resistivity of the ionized and turbulent plasma. Then a driven pulsed jet or a collision with the disk by a star, Fig.~\\ref{flux_schem}(b(C))] causes a plume to rise either towards the end plate or above the disk with the corresponding displaced toroidal flux forming a loop of toroidal flux. The radial expansion of the plume material causes the plane of this loop to untwist or rotate differentially about its own axis relative to the rotating frame so that the initial toroidal orientation of the loop is transformed to a poloidal one, Fig.~\\ref{flux_schem}(b(D)). Resistive diffusion in liquid sodium metal or reconnection in the ionized plas- \\\\ ma of the disk allows this now poloidal loop to merge with the original poloidal field. For positive dynamo gain, the rate of addition of poloidal flux must be greater than its decay. It is only because the toroidal multiplication can be so large or that $R_{m,\\Omega}$ can be so large that the helicity necessary for gain of the $\\alpha\\omega$ dynamo can be much smaller and episodic. ", + "conclusions": "The numerical simulations along with the engineering design feasibility give us confidence that, with a continuing effort a positive gain $\\alpha\\omega$ dynamo can be made in the laboratory." + }, + "0112/astro-ph0112407_arXiv.txt": { + "abstract": "\\noindent We have found a very faint companion to the active solar analog HR~7672 (HD~190406; GJ~779; 15 Sge). Three epochs of high resolution imaging using adaptive optics (AO) at the Gemini-North and Keck~II Telescopes demonstrate that HR~7672B is a common proper motion companion, with a separation of 0\\farcs79 (14~AU) and a 2.16~\\micron\\ flux ratio of 8.6~mags. Using follow-up $K$-band spectroscopy from Keck AO+NIRSPEC, we measure a spectral type of L4.5$\\pm$1.5. This is the closest ultracool companion around a main sequence star found to date by direct imaging. We estimate the primary has an age of 1--3~Gyr. Assuming coevality, the companion is most likely substellar, with a mass of 55--78~\\Mjup\\ based on theoretical models. The primary star shows a long-term radial velocity trend, and we combine the radial velocity data and AO imaging to set a firm (model-independent) lower limit of 48~\\Mjup. In contrast to the paucity of brown dwarf companions at $\\lesssim$4~AU around FGK dwarfs, HR~7672B implies that brown dwarf companions do exist at separations comparable to those of the giant planets in our own solar system. Its presence is at variance with scenarios where brown dwarfs form as ejected stellar embryos. Moreover, since HR~7672B is likely too massive to have formed in a circumstellar disk as planets are believed to, its discovery suggests that a diversity of physical processes act to populate the outer regions of exoplanetary systems. ", + "introduction": "Radial velocity surveys find that about 6\\% of solar-type FGK stars harbor planets within 4~AU, with $M \\sin i$ spanning 0.25--15~\\Mjup. In contrast, the same surveys find the incidence of more massive substellar companions (15--80~\\Mjup) at these radii is $\\lesssim$1\\%, even though such objects are much easier to detect. The evidence for this ``brown dwarf desert'' first emerged from radial-velocity surveys in the late 1980's and early 90's \\citep{1988ApJ...331..902C, 1989ApJ...344..441M, 1995Icar..116..359W}. Their modest precisions ($\\sim$300~m/s) were sufficient to detect brown dwarfs inside 3~AU, but only a single candidate was found.\\footnote{This object was the companion to HD~114762 with $\\Msini = 11~\\Mjup$ \\citep{1989Natur.339...38L}. The primary star is likely seen pole-on so its companion was believed to be stellar \\citep{1991ApJ...380L..35C}. Recent AO imaging has resolved a third component, a late-type dwarf at 91~AU \\citep{lloyd00}. This raises the possibility of dynamical interactions causing the stellar rotation axis to be mis-aligned with the companion orbital axis \\citep{pat01}. Hence, HD~114762B may in fact be a substellar companion.} Current high-precision ($\\sim$10~m/s) radial velocity programs have identified about a dozen close companions with $M \\sin i = 15-60~\\Mjup$ \\citep{1997abos.conf..313M}. However, the majority of these have been found by {\\sl Hipparcos} to be unresolved astrometric binaries with M-dwarf companions, and the remainder are unlikely to be brown dwarfs \\citep{2000A&A...355..581H, zuck01, 2001A&A...372..935P}. This paucity of objects stands in stark contrast to the abundance of free-floating brown dwarfs in the field \\citep[e.g.][]{1999ApJ...521..613R, 2000AJ....120..447K} and in young clusters down to very low masses \\citep[e.g.][]{1998A&A...336..490B, 2000ApJ...540.1016L, 2000ApJ...541..977N, liu01a}. An unanswered observational question is whether the brown dwarf desert exists at $\\gtrsim$4~AU, outside the region of current radial velocity surveys. A few L and T-dwarf companions have been found around FGK dwarfs from the 2MASS imaging survey at very wide separations (250--2500~AU) \\citep{2000ApJ...531L..57B, 2001AJ....121.3235K, wils01}. The statistics are currently small, but this suggests the brown dwarf desert does not exists at very large separations \\citep{2001ApJ...551L.163G}. These companions might have formed during the fragmentation of the same natal molecular core as the primary. Given their mass and large separation, it is unlikely that they formed in a circumstellar disk, as planets are believed to form. Little is known about the frequency of substellar companions at $\\approx$4--50~AU. In our solar system, this is the domain of giant planets and the Kuiper Belt. Hence, probing these separations around other stars can test our understanding of formation processes in the outer regions of planetary systems. There are clues that massive planets and/or brown dwarfs do exist at these radii: half of stars with planets have systematic long-term trends in their radial velocities due to unseen companions \\citep{2001ApJ...551.1107F}. The orbital periods are much longer than the observing baselines, and hence the masses are poorly constrained. Because the periods are many years and/or decades, relying on radial velocities alone will require a long time to determine the companion masses. However, adaptive optics (AO) imaging with large ground-based telescopes can probe the physical nature of these companions and lend insight into the mass distribution of substellar and planetary companions. Thus, ground-based AO imaging provides a key capability for finding substellar companions to main sequence stars, providing sensitivity at radii outside radial velocity searches but closer than ordinary imaging surveys. Recently, \\citet{els01} have found a common proper motion companion to the extrasolar planet star GL~86. Their coronagraphic AO imaging finds a companion at a separation of 1\\farcs7 (19~AU). The companion is among the coolest known, probably at the transition between L and T-type objects based on IR photometry. Aside from implications for understanding the planetary formation process, brown dwarf companions to main sequence stars are interesting in their own right. Since a brown dwarf does not have a stable source of internal energy, its age and mass are degenerate for a given spectral type --- young lower mass BDs can have same temperature as older higher mass BDs. By finding brown dwarf companions to main sequence stars, we can break this degeneracy by measuring the age of the primary and assuming the components are coeval. Combined with theoretical models for the cooling history of substellar objects, we can determine masses for brown dwarfs. Here we report the discovery and characterization of a common proper motion companion to HR~7672 (HD~190406, 15~Sge, GJ~779) using the AO systems at the Gemini-North and Keck~II Telescopes. The star lies at a distance of 17.7~pc and has a spectral type of G1V. \\citet{1996A&ARv...7..243C} classified it as a solar analog, and its higher level of activity than the Sun suggests a young age. Our observations and astrometry of HR~7672B appear in \\S~2. We present temperature and mass determinations for the companion based on near-IR AO photometry, spectroscopy, and radial velocity data in \\S~3. In \\S~4, we summarize our results and offer some implications for the formation of substellar objects around solar-type stars. ", + "conclusions": "We have found a common proper motion companion to the nearby solar analog HR~7672 using AO imaging from the Gemini and Keck Telescopes. IR spectra from Keck AO+NIRSPEC find the companion has a very cool atmosphere, and its faint $K$-band magnitude also points to a very late-type object: we estimate a spectral type of L4.5$\\pm$1.5. This alone suggests the companion is sustellar, or right at the substellar boundary. A variety of data indicate the primary star is younger than the Sun but older than the Hyades, with a likely age around 1--3~Gyr. Using theoretical models for cooling of very low mass objects, the inferred companion mass ranges from 55--78~\\Mjup, under the assumption that the two components are coeval. The primary has a long-term radial velocity acceleration, and we combine the radial velocity and AO data to set a lower limit of 48~\\Mjup\\ on the companion mass strictly from dynamical considerations. At a separation of only 0\\farcs79 (14~AU), this is the closest ultracool (L or T-dwarf) companion to a main sequence star found to date by direct imaging, in both angular and physical separation (see compilation in \\citealp{2001AJ....121..489R}). For a random distribution of orbital eccentricities, $\\approx$85\\% of orbits have a true semi-major axis which is 0.5--2~times that of the observed separation. Hence the semi-major axis of HR~7672B is likely to be 7--28~AU, meaning an orbital period of 20--150~yrs. Orbital motion could be detectable in a few years. This raises the possibility of determining its orbital eccentricity, perhaps a important clue into its formation history. Furthermore, the combination of AO imaging with continued radial velocity monitoring can be used to better constrain the companion mass in advance of observing a full orbital period. The formation and presence of close brown dwarf companions might inhibit the formation of circumstellar disks, and hence planets. Using radial velocity data, \\citet{1999ApJ...526..890C} set upper limits on any planetary companion of 0.1--10~\\Mjup\\ at separations of 0.1--4~AU, respectively. Hence, we know the HR~7672 system does not have any Jovian mass planets in its inner regions. However, we know from the case of GL~86 that at least one brown dwarf co-exists with an extrasolar planet \\citep{els01}. The existence of HR~7672B and GL~86B indicates that brown dwarfs can form at separations comparable to circumstellar disks. (These objects are most likely substellar, or right at the substellar boundary. In the discussion which follows, we assume for the sake of arugment that these are brown dwarfs.) A third possible example is the brown dwarf companion orbiting HD~168443 \\citep{2001ApJ...555..418M}, at a semi-major axis of 3~AU and with a mass between 17.2~\\Mjup\\ (from the radial velocities) and 42~\\Mjup\\ (from limits on any astrometric wobble). This system also illustrates the gradual convergence of radial velocities and AO imaging, closing the mass gap for discovery of planets and brown dwarfs by indirect and direct means. Formation scenarios for HR~7672B, and their connection with ones for extrasolar planets, remain an open question. Current semantic convention designates planets as object below $\\sim$15~\\Mjup, with this border motivated either by the deuterium-burning limit for substellar objects or simply by the observed steep decline in frequency of such objects in radial velocity surveys \\citep{2001ApJ...555..418M}; more massive objects are considered brown dwarfs. HR~7672B would certainly be considered in the brown dwarf regime based on its estimated mass. \\citet{2001AJ....122..432R} have suggested brown dwarfs form as stellar embryos in small newborn multiple systems: they grow in mass by accretion of infalling gas, but their growth is prematurely truncated by ejection due to dynamical interactions. The discovery of HR~7672B is at variance with this scenario, since the ejection model discounts the existence of close separation brown dwarfs around solar-type stars. An additional criterion to distinguish planets from brown dwarfs under consideration is whether the object forms ``as planets do,'' presumably in a circumstellar disk, or instead as an isolated object. In this regard, objects like HR~7672B and GL~86B are interesting in that they resides in the zone of giant planet formation, \\ie, $\\sim$~5--30~AU where we know giant planets exist in our own solar system and theoretical models can most easily form them. Giant planets are thought to form at these radii via gas accretion onto a rocky core. Subsequent dynamical processes may cause them to migrate to smaller radii \\citep{1996Natur.380..606L, 1997ApJ...482L.211W, 1998ApJ...500..428T}, since the $\\lesssim$4~AU planets found by radial velocity surveys are not thought to have formed in situ. However, brown dwarfs are believed to be too massive to form by the same core accretion scenario as giant planets. Instead, they may have originated, \\eg, by fragmentation during initial cloud collapse \\citep{1998bdep.conf..115B}, by instabilities in a very massive disk \\citep{1998ApJ...503..923B, 2000ApJ...540.1091B, 2000ApJ...540L..95P}, or by collisions between protoplanetary disks \\citep{1998Sci...281.2025L}. Furthermore, brown dwarf companions are not found at small separations, but we now have strong evidence that they exist at $\\sim$15~AU. This suggest they are immune to the migration process(es) which affect massive planets. The discovery of brown dwarfs in the zone of giant planet formation implies that a diversity of physical processes act to populate the outer regions of exoplanetary systems." + }, + "0112/hep-ph0112238_arXiv.txt": { + "abstract": "Phenomenology of relic neutralinos is analyzed in an effective supersymmetric scheme at the electroweak scale. It is shown that current direct experiments for WIMPs, when interpreted in terms of relic neutralinos, are indeed probing regions of supersymmetric parameter space compatible with all present experimental bounds. ", + "introduction": "\\noindent Interest for supersymmetric relics leans on the following properties: i) if R--parity is conserved, the Lightest Supersymmetric Particle (LSP) is stable; ii) if colourless and uncharged, the LSP is a nice realization of a relic Weakly Interacting Massive Particle (WIMP) \\cite{bf}. In the present paper we assume that supersymmetry exists in Nature and that properties (i) and (ii) hold. Thus the LSP will be our relic WIMP candidate; actually, after some preliminary considerations, our attention will be mainly devoted to the case when the LSP is provided by the neutralino. In Sect. 2 we discuss some generic connections between WIMP relic abundance and event rates for direct and indirect WIMP detection. Sect. 3 is devoted to a presentation of supersymmetric schemes. Numerical results and conclusions are finally presented in Sect. 4. ", + "conclusions": "In what follows we present some of our results in the framework of the effMSSM scheme. The evaluation of $\\Omega_{\\chi} h^2$ follows the procedure given in \\cite{noiom}. The cross--section $\\sigma^{\\rm (nucleon)}_{\\rm scalar}$ has been calculated with the formulae reported in Refs.\\cite{noi6}. Important entries in $\\sigma^{\\rm (nucleon)}_{\\rm scalar}$ are the quantities $m_q $, where the quark scalar densities $\\bar q q$ are averaged over the nucleonic state. The values of these quantities, derived from the pion--nucleon sigma term $\\sigma_{\\pi N}$ and other hadronic quantities, are affected by large uncertainties \\cite{noi6}. The quantity $m_s $, which is the most important term among the $m_q $'s unless $\\tan \\beta$ is very small \\cite{ggr}, is affected by an uncertainty factor larger than 3. This conclusion is reinforced by the most recent determinations of $\\sigma_{\\pi N}$ \\cite{olsson,pavan}, as discussed in Ref. \\cite{size}. The results presented in Figs. 1--3 employ the following set of values for the quantities $m_q $ (denoted as set 1 in Ref. \\cite{noi6}): $ m_{l}<\\bar{l}l>\\; =\\; 23\\; {\\rm MeV}, \\; m_{s}<\\bar{s}s>\\; =\\; 215\\; {\\rm MeV}, \\; m_{h}<\\bar{h}h>\\; =\\; 50\\; {\\rm MeV}$. For the reasons discussed above, all values concerning $\\sigma_{\\rm scalar}^{(\\rm nucleon)}$ in Figs. 1--3 are subject to an increase of $\\sim$ 4, when the current uncertainties in the $m_q $'s are taken into account. \\begin{figure}[htb] \\centering \\includegraphics[height=4in]{cosmo_fig1_lr.ps} \\caption{Scatter plot of $\\sigma_{\\rm scalar}^{(\\rm nucleon)}$ versus $\\Omega_{\\chi} h^2$. $m_{\\chi}$ is taken in the the interval $40 \\; {\\rm GeV} \\leq m_W \\leq 200 \\; {\\rm GeV}$. The two horizontal lines bracket the sensitivity in current WIMP direct experiments $4 \\cdot 10^{-10} \\; {\\rm nbarn} \\leq \\ \\xi \\sigma^{\\rm (nucleon)}_{\\rm scalar} \\leq 2 \\cdot 10^{-8} \\; {\\rm nbarn}$, for $\\xi = 1$. The two vertical lines denote the range $0.05 \\leq \\Omega_{m} h^2 \\leq 0.3$. Dots (crosses) denote gaugino (mixed) configurations. } \\end{figure} Now let us turn to the presentation of our main results. In Fig.1 we give the scatter plot for $\\sigma_{\\rm scalar}^{(\\rm nucleon)}$ versus $\\Omega_{\\chi} h^2$. The two vertical lines denote the favorite range for $\\Omega_{m}h^2$: $0.05 \\leq \\Omega_{m} h^2 \\leq 0.3$. The two horizontal lines bracket the range of sensitivity in current WIMP direct experiments \\cite{ge,dama}, which, taking into account astrophysical uncertainties \\cite{belli}, turns out to be $4 \\cdot 10^{-10} \\; {\\rm nbarn} \\leq \\ \\xi \\sigma^{\\rm (nucleon)}_{\\rm scalar} \\leq 2 \\cdot 10^{-8} \\; {\\rm nbarn}$, for WIMP masses in the interval $40 \\; {\\rm GeV} \\leq m_W \\leq 200 \\; {\\rm GeV}$. Fig.1 shows that the present experimental sensitivity in WIMP direct searches allows the exploration of supersymmetric configurations compatible with current accelerator bounds; a number of configurations stay inside the region of cosmological interest. \\begin{figure}[htb] \\centering \\includegraphics[height=4in]{cosmo_fig2_lr.ps} \\caption{ Scatter plot of $\\rho_{\\chi}$ versus $\\Omega_{\\chi}h^2$. This plot is derived from the experimental value $[\\rho_{\\chi}$/(0.3 GeV cm$^{-3}$) $\\cdot \\sigma^{\\rm (nucleon)}_{\\rm scalar}]_{expt} = 1 \\cdot 10^{-9}$ nbarn and by taking $m_{\\chi}$ in the the interval $40 \\; {\\rm GeV} \\leq m_W \\leq 200 \\; {\\rm GeV}$, according to the procedure outlined in the text. The two horizontal lines delimit the range 0.2 GeV cm$^{-3} \\leq \\rho_{\\chi} \\leq $ 0.7 GeV cm$^{-3}$; the two vertical ones delimit the range $0.05 \\leq \\Omega_{m} h^2 \\leq 0.3$. The band delimited by the two slanted dot--dashed lines and simply hatched is the region where rescaling of $\\rho_l$ applies. Dots denote gauginos, circles denote higgsinos and crosses denote mixed configurations. } \\end{figure} Once a measurement of the quantity $\\rho_{\\chi} \\cdot \\sigma^{\\rm (nucleon)}_{\\rm scalar}$ is performed, values for the local density $\\rho_{\\chi}$ versus the relic abundance $\\Omega_{\\chi}h^2$ may be deduced by proceeding in the following way \\cite{noi6}: 1) $\\rho_{\\chi}$ is evaluated as $[\\rho_{\\chi} \\cdot \\sigma^{\\rm (nucleon)}_{\\rm scalar}]_{expt}$ / $\\sigma^{\\rm (nucleon)}_{\\rm scalar}$, where $[\\rho_{\\chi} \\cdot \\sigma^{\\rm (nucleon)}_{\\rm scalar}]_{expt}$ denotes the experimental value, and $\\sigma^{\\rm (nucleon)}_{\\rm scalar}$ is calculated as indicated above; 2) to each value of $\\rho_{\\chi}$ one associates the corresponding calculated value of $\\Omega_{\\chi} h^2$. The scatter plot in Fig.2 is derived for the representative value $[\\rho_{\\chi}$/(0.3 GeV cm$^{-3}$) $\\cdot \\sigma^{\\rm (nucleon)}_{\\rm scalar}]_{expt} = 1 \\cdot 10^{-9}$ nbarn within the annual--modulation region of Ref. \\cite{dama}, and by taking $m_{\\chi}$ in the range $40 \\; {\\rm GeV} \\leq m_W \\leq 200 \\; {\\rm GeV}$. The plot of Fig.2 shows that the most interesting region, {\\it i.e.} the one with 0.2 GeV cm$^{-3} \\leq \\rho_{\\chi} \\leq $ 0.7 GeV cm$^{-3}$ and $0.05 \\leq \\Omega_{m} h^2 \\leq 0.3$ (cross-hatched region in the figure), is covered by susy configurations probed by the WIMP direct detection. Let us examine the various sectors of Fig.2. Configurations above the upper horizontal line are incompatible with the upper limit on the local density of dark matter in our Galaxy and must be disregarded. Configurations above the upper slanted dot--dashed line and below the upper horizontal solid line would imply a stronger clustering of neutralinos in our halo as compared to their average distribution in the Universe. This situation may be considered unlikely, since in this case neutralinos could fulfill the experimental range for $\\rho_\\chi$, but they would contribute only a small fraction to the cosmological cold dark matter content. For configurations which fall inside the band delimited by the slanted dot--dashed lines and simply--hatched in the figure, the neutralino would provide only a fraction of the cold dark matter at the level of local density and of the average relic abundance, a situation which would be possible, for instance, if the neutralino is not the unique cold dark matter particle component. To neutralinos belonging to these configurations one should assign a {\\it rescaled} local density. We remind that the scatter plot in Fig.2 refers to a representative value of $[\\rho_{\\chi}$ $\\cdot \\sigma^{\\rm (nucleon)}_{\\rm scalar}]$ inside the current experimental sensitivity region, thus the plot in Fig.2 shows that current experiments of WIMP direct detection are probing relic neutralinos which may reach values of cosmological interest, but also neutralinos whose local and cosmological densities may provide only a very small fraction of these densities. \\begin{figure}[htb] \\centering \\includegraphics[height=4in]{cosmo_fig3_lr.ps} \\caption{Scatter plot of $\\xi \\sigma_{\\rm scalar}^{(\\rm nucleon)}$ versus $m_{\\chi}$. Crosses (dots) denote configurations with $\\Omega_{\\chi} h^2 > 0.05$ ($\\Omega_{\\chi} h^2 < 0.05$). The solid contour denotes the 3$\\sigma$ annual--modulation region of Ref.\\cite{dama} (with the specifications given in the text). } \\end{figure} In Fig.3 we give the scatter plot for the quantity $\\xi \\sigma^{\\rm (nucleon)}_{\\rm scalar}$ versus $m_{\\chi}$. The solid line denotes the frontier of the 3$\\sigma$ annual--modulation region of Ref.\\cite{dama}, when only the uncertainties in $\\rho_l$ and in the dispersion velocity of a Maxwell--Boltzmann distribution, but not the ones in other astrophysical quantities, are taken into account. Effects due to a possible bulk rotation of the dark halo or to an asymmetry in the WIMP velocity distribution would move this boundary towards higher values of $m_{\\chi}$ \\cite{belli}. Our results in Fig.3 show that the susy scatter plot reaches up the annual--modulation region of Ref.\\cite{dama}. In our figures only results referring to the effMSSM scheme are reported. For comparisons among results in various schemes: universal SUGRA, non--universal SUGRA and effMSSM, we refer to Refs. \\cite{lathuile,probing}. As for the universal SUGRA, we only wish to remark that this very constrained model, combined with the present rather stringent experimental bounds from LEP2, typically entails a sizeable suppression of the neutralino--nucleon cross--section. Whether or not this suppression may prevent the calculated $\\sigma_{\\rm scalar}^{(\\rm nucleon)}$ from reaching the region of present experimental sensitivity does depend on how the various constraints (typically the bounds on Higgs masses, on $m_t$ and $m_b$) are implemented in the evaluations. By way of example, it is worth mentioning that the explicit bounds on the quantity $\\sin^2 (\\alpha - \\beta)$ ($\\alpha$ being the Higgs mixing angle in the neutral CP--even Higgs sector) as a function of $m_h$ should be taken into account, rather than using a flat lower bound of 115 GeV for $m_h$ \\cite{higgslep,w}. Neglecting these features in a SUGRA calculation may lead to biased conclusions. We now summarize the main points of this paper: \\begin{itemize} \\item Most recent theoretical developments suggest supersymmetric schemes which notably differ from a strict model such as the universal SUGRA and point to the fact that this constrained scheme should be relaxed in many instances. Here, we have employed an effective MSSM scheme at the electroweak scale, which is particularly convenient to treat the relic neutralino phenomenology. \\item We have shown that current direct experiments for WIMPs, when interpreted in terms of relic neutralinos, are indeed probing regions of supersymmetric parameter space compatible with all present experimental bounds. \\item We have proved that part of the configurations probed by current WIMP experiments entail relic neutralinos of cosmological interest, and, {\\it a fortiori} also neutralinos which might comprise only a fraction of the required amount of dark matter in the Universe. \\end{itemize}" + }, + "0112/astro-ph0112361_arXiv.txt": { + "abstract": "Redshift distortions, both geometrical and kinematical, of quasar clustering are simulated, for the Two-Degree Field QSO Redshift Survey (2QZ), showing that they are very effective to constrain the cosmological density and equation of state parameters, $\\Omega_{m0}$, $\\Omega_{x0}$ and $w$. Particularly, it emerges that, for the cosmological constant case, the test is especially sensitive to the difference $\\Omega_{m0}-\\Omega_{\\Lambda 0}$, whereas, for the spatially flat case, it is quite competitive with future supernova and galaxy count tests, besides being complimentary to them. ", + "introduction": "Following Alcock \\& Paczy\\'nski's (1979) and Phillipps' (1994) lead, we extended Popowski et al. (1998) investigation as follows: (1) We performed Monte Carlo simulations to obtain the probability density function and corresponding confidence contours in the parametric plane $(\\Omega_{m0},\\Omega_{\\Lambda 0})$, comparing them to other tests; (2) We included a general dark energy component with constant equation of state parameter $w$, obtaining, for flat models, the confidence contours in the $(\\Omega_{m0},w)$ plane; (3) We explicitly took into account the effect of large-scale coherent peculiar velocities (Hamilton 1992; Matsubara \\& Suto 1996). Our calculations are based on the measured 2QZ distribution function and we consider best fit values for the amplitude and exponent of the correlation function as obtained by Croom et al. (2001). ", + "conclusions": "Fig.~1 shows that the test is quite sensitive (when $w=-1$) to the difference $\\Delta:=\\Omega_{m0}-\\Omega_{\\Lambda 0}$; also, for flat models, the constraints are similar to those from SNAP and somewhat better than the ones from DEEP. We have also run simulations where the ``true'' model takes into account the linear kinematical redshift distortion and the simulated ones do not; it then turns out that the ``true'' model is not reliably recovered (at $2\\sigma$ level), demonstrating the importance of taking this effect explicitly into account. In contrast, for different bias functions for the ``true'' and the simulated models, we were able to faithfully recover $\\Delta$. We conclude by stressing that this sort of Alcock-Paczy\\'nski test is very promising and referring the reader to a more detailed version of this work (Calv\\~ao, de Mello Neto \\& Waga 2001)." + }, + "0112/astro-ph0112157_arXiv.txt": { + "abstract": "We present high-resolution 1.4 GHz Australia Telescope Compact Array polarimetric observations of Vela X, the pulsar wind nebula of the Vela SNR. We find that the linearly polarized emission is only partially correlated with total intensity. There are many depolarization features, some of which are coincident with foreground H$\\alpha$ filaments. Further study of these should provide measurements of the magnetic field in the remnant's shell. ", + "introduction": "Vela X is the pulsar wind nebula near the center of the Vela supernova remnant (Bock, Turtle, \\& Green 1998b). In total intensity, the radio emission is dominated by synchrotron filaments which have no clear optical counterpart (Bock et al.\\ 1998a). The remnant's distance of 350 pc\\footnote{Recent measurements of the pulsar's distance include an optical parallax of $294_{-50}^{+76}$ pc (Caraveo et al.\\ 1991) and a radio parallax of $410\\pm42$ pc (Legge 2001)} makes it one of the nearest SNRs, and thus one of the easiest to study. Its age, 10,000 yr, is much larger than that of the Crab Nebula, and we can hope to learn something about remnant evolution by comparing the sources. High quality interferometric polarimetry now makes it possible to study at sub-arcminute resolution the magnetic fields in pulsar-powered nebulae, and in the Galaxy as a whole. ", + "conclusions": "\\begin{figure} \\centering \\includegraphics[height=11cm]{bockd2_2.eps} \\caption{Polarization E-vectors overlaid on total intensity. The peak total intensity of the extended emission is 220 mJy~beam$^{-1}$. The synthesized beam and a bar representing the maximum polarized intensity of 106 mJy~beam$^{-1}$ are shown at the lower right. The three unresolved sources are most likely of extragalactic origin. } \\label{fig:ppa} \\end{figure} The most striking aspect of the polarimetric image of Vela X (Figure \\ref{fig:p}) is the complex network of `canals' criss-crossing the broader synchrotron filaments. These canals are unresolved, indicating that they are probably due to Faraday depolarization within the synthesized beam, caused by rapid spatial changes in the foreground rotation measure (RM). They have no counterpart in total intensity. Similar features have been seen in the background Galactic radiation (Gaensler et al.\\ 2001, and references therein). However, in the present case we have been able to identify optical counterparts for many of the canals among H$\\alpha$ filaments originating in the shell of the Vela SNR (Figure \\ref{fig:ha}). By considering the geometry of the filaments (which are probably sheets in projection) and obtaining electron densities in the region from further optical studies, we expect to be able to make a \\emph{direct} measurement of the compressed magnetic field in the Vela SNR shell. The underlying more diffuse linearly polarized emission has some overall corellation with total intensity. However, there are many regions of disagreement. These could be either intrinsic to the source (i.e.\\ due to variations within the internal magnetic fields) or result from depolarization by some intervening more compact region. In the latter case, their counterpart in total intensity is the underlying diffuse emission which forms the majority of the flux density measured from Vela X. The polarization E-vector direction (Figure~\\ref{fig:ppa}) is a useful diagnostic of the magnetic field. In this region, the RM is approximately 40 rad~m$^{-2}$ (Milne 1995), indicating that the magnetic field is generally aligned with the E-vectors plotted, and thus lies along the filaments (the discrepant E-vectors half-way along the central filament correspond to a region of higher RM). We note that the `wisp' just to the north-east of the pulsar (Bietenholz, Frail, \\& Hankins 1991) and the filament extending south do not appear symmetric across the pulsar. Further work to make high-resolution corrections for Faraday rotation will allow a more detailed interpretation." + }, + "0112/astro-ph0112227_arXiv.txt": { + "abstract": "Recently the AGASA Collaboration presented data suggesting a significant clustering of ultra-high energy cosmic rays coming from the outer Galaxy region. In this paper we calculate expected cosmic ray arrival distributions for several simple, limiting source location scenarios and investigate the possibility of clustering and correlation effects. The role of the Galactic magnetic field is discussed in detail. ", + "introduction": "The observations of cosmic rays of energy above 10$^{19}$ eV reveal at least two features that we have not yet understood and that appear to contradict each other. The cosmic ray energy spectrum does not seem to be cut off because of photoproduction interactions on the microwave background \\cite{Greisen:1966,Zatsepin:1966} and extends above 10$^{20}$ eV \\cite{Takeda:1998ps}. The existence of these particles suggests that the sources of ultra-high energy cosmic rays (UHECRs) are cosmologically nearby - within 20 Mpc or so \\cite{Stanev:2000fb}. On the other hand, the UHECR arrival direction distribution, as far as we know it, seems isotropic on a large scale with a statistically improbable small scale clustering \\cite{Uchihori:1999gu} (for a review on the subject see \\cite{Nagano:2000ve}). This small-scale clustering was first observed in the data of the AGASA group \\cite{Takeda:1999ap}. The current AGASA data set of 59 events of energy above $4\\times10^{19}$ eV contains five doublets and one triplet for a maximum separation angle of $2.5^\\circ$~\\cite{Takeda:2001icrc}. The centroids of the triplet and one of the doublets are less than 1 degree off the supergalactic plane (SGP). When combined with data from the Haverah Park, Yakutsk and Volcano Ranch experiments the number of doublets (within $3^\\circ$) increases to eight and the number of triplets to two \\cite{Uchihori:1999gu}. The chance probability of observing these multiplets from a uniform distribution of sources is less than $1.5~\\%$ and less than $1~\\%$ when a restricted region within $\\pm10^{\\circ}$ off the SGP is considered \\cite{Uchihori:1999gu}. Recently the AGASA group presented a self-correlation analysis of their data with $E>10^{19}$ eV \\cite{Takeda:2001icrc}. Remarkable correlations are found in Galactic coordinates, supporting further the previously found indications of clustering of UHECRs. In a straight-forward interpretation one would link the existence of such clusters directly to UHECR point sources \\cite{Tinyakov:2001ic,Anchordoqui:2001qk} but non-trivial effects such as possible clustering of sources or large-scale magnetic fields could also contribute to such a correlation \\cite{Stanev:1996qj,MedinaTanco:1998aj,% MedinaTanco:1998ap,MedinaTanco:1998yq,MedinaTanco:1999aj,% Harari:2000he, Harari:2000az}. The indications found in data raise the interesting question of whether both an isotropic distribution and a clustering of the arrival directions can be explained consistently, assuming they are not a statistical fluctuation \\cite{Goldberg:2000zq,Razzaque:2001tp}. Potential models of UHECR origin, assuming they are charged particles, are constrained by their ability to reproduce the measured energy spectrum and the approximately isotropic arrival distribution observed in the data \\cite{Olinto:2000sa,Bhattacharjee:1998qc}. At the same time, their predictions should also be consistent with the small scale clustering. In general, several key ingredients enter the models: (i) the distribution (locations) of the sources, (ii) their nature, i.e. whether they are sources emitting cosmic rays continuously for a long period of time or they are bursting sources, (iii) the energy spectrum and total flux at injection, and (iv) the propagation of the cosmic rays including energy loss processes and deflection due to magnetic fields. Clustering on small scales and overall isotropy seems to favour models based on several nearby sources (for example, in the Galactic halo). Models of this type include the acceleration of cosmic rays at magnetars \\cite{Blasi:1999xm}, UHE neutrino interactions on the relic neutrino background (Z-bursts)~\\cite{Weiler:1997sh,Fargion:1997ft}, and the decay and annihilation of superheavy Relics accumulated in the halo of the Galaxy \\cite{Berezinsky:1997hy,Berezinsky:1998ed,Birkel:1998nx,% Blasi:2001hr,Evans:2001rv,Sarkar:2001se}. In this paper we study general aspects related to cosmic ray propagation and the source distribution. The aim is to explore what a future confirmation of the correlations could reveal about the nature of the sources of UHECRs, their distribution and the strength of the magnetic field in our Galaxy. To keep our results as general as possible we consider three simplified, limiting source location scenarios:\\\\ (i) Uniform distribution of sources, which is relevant to UHECR of extragalactic origin with sources of isotropic and homogeneous distribution.\\\\ (ii) sources distributed uniformly within $\\pm20^\\circ$ off the supergalactic plane (SGP), and\\\\ (iii) single and multiple point sources.\\\\ For each of these source scenarios we calculate arrival distributions, one and two-dimensional correlation functions and, where applicable, arrival time delays. It should be emphasized that we do not attempt to reproduce AGASA data with these source models. The interest of this work is merely the investigation of characteristic features which are fairly model-independent. For a more realistic analysis one needs a better knowledge of the Galactic magnetic field, as well as full access to the experimental data, including acceptance corrections. The outline of the paper is as follows. In Sec.~\\ref{sec:mag-field} we discuss the Galactic magnetic field. The simulation method for cosmic ray arrival distributions is explained in Sec.~\\ref{sec:method} and the different source location scenarios are presented in Secs.~\\ref{sec:uniform} through \\ref{sec:multiple-point}. Sec.~\\ref{sec:discussion} applies the results of the previous sections to AGASA multiplets and Sec.~\\ref{sec:end} concludes the paper with a summary of our main results. ", + "conclusions": "} We have performed a detailed calculation of the UHECR flux expected in several simple, limiting source scenarios. We have implemented an approach different from that used by other authors propagating the particles from the sources to the detector. This allows not only the investigation of the particle's deflection in the GMF but also the detailed study of various source scenarios. We find the self--correlation analyses in one and two dimensions to be powerful tools for identification of the sources of UHECR and their distribution. The predictive power of our results is mainly limited by the uncertainty in the GMF model we use in terms of structure, field strength and spatial extent. Nevertheless we have obtained a number of interesting results which do not depend very much on the parameters of the particular GMF model. \\begin{itemize} \\item For an isotropic extragalactic flux one expects an isotropic flux at Earth as long as cosmic rays are not trapped by the GMF. The observed CRs would stem from sources being distributed non-uniformly. \\item The SGP plane is disfavoured as main source of UHECR as long as the extragalactic magnetic field is weak. \\item Point sources produce two-dimensional correlation ellipses, the slope of which depends on the position of the source and not on its distance. \\item The observed cosmic ray flux of point sources is energy-dependent modulated. For example, cosmic rays with $10^{19}$ eV are suppressed from sources in the Galactic plane. \\item The size of the correlation in $\\Delta l - \\Delta b$ depends strongly on the strength of the GMF and only weakly on the distance to the sources. \\item The energy-ordered correlation analysis reveals the general field direction/polarity in the vicinity of the Solar system. It can be used to determine the charge sign of the cosmic rays and also helps to estimate the importance of statistical fluctuations in data. \\item The width of the time delay versus energy correlation allows for the inverted arrival time behaviour, i.e. the higher energy cosmic ray of a pair arrives earlier than the lower energy one. This may happen if the energies of the cosmic rays are within $\\sim 30\\%$ of each other. \\item It seems conceivable that a scenario of multiple, uniformly distributed sources can produce correlation patterns as seen in AGASA data. \\item The slope and the lobes of the 2D correlation ellipses are not statistically significant for less than $\\sim 200$ events from the outer Galaxy in the multiple sources scenario. \\item The energy-ordered 2D correlation distribution of the AGASA data does not agree with the expectations for positively charged particles and the BSS GMF model, however the small statistics of the data does not allow us to draw any firm conclusions. \\item The analysis of the AGASA multiplets within a model of point sources shows that only one doublet is excluded from being a pair of cosmic rays simultaneously emitted from a source at a distance between 15 - 40 kpc. All other multiplets are within the uncertainties compatible with coming from short-lived point sources. \\item Further three possible multiplets were selected on the basis of the expected self-correlation patterns. \\end{itemize} In forthcoming work we plan to derive limits on UHECR models such as $Z$-bursts and decay of super-heavy particles. Supposing the self-correlation observed in AGASA data is not a statistical fluctuation, we will extend our work to set limits on the minimum number of nucleon-antinucleon pairs needed per point source to produce doublets and triplets as well as limits on the total number of point sources required to sustain the observed UHECR flux. \\vspace*{1cm} \\noindent {\\bf Acknowledgements} The authors acknowledge fruitful discussions with P.L. Biermann, T.K. Gaisser, P.P. Kronberg and M. Teshima. This research is supported in part by NASA Grant NAG5-7009. RE\\&TS are also supported by the US Department of Energy contract DE-FG02 91ER 40626. TS also acknowledges the hospitality of PCC, Coll\\`ege de France and a grant from the Minist\\`ere de la Recherche of France. The simulations presented here were performed on DEC Alpha and Beowulf clusters funded by NSF grant ATM-9977692." + }, + "0112/astro-ph0112005_arXiv.txt": { + "abstract": "There is growing evidence that some variable high-energy $\\gamma-$ray sources are associated with pulsar wind nebulae. We review the current status of X-ray and radio studies of the sources most likely to be significant emitters above 100 MeV. The assumption of the $\\gamma-$radiation arising from synchrotron processes puts a lower limit on the magnetic field and suggests there may be significant doppler boosting causing the asymmetries seen in the morphology of these nebulae. However, these sources are also in close proximity to molecular clouds, suggesting that variable $\\gamma-$ray emission may arise from a ram pressure confined pulsar wind nebula in a dense medium. ", + "introduction": "For over 30 years, the nature of the majority of high-energy ($E > 100$ MeV) $\\gamma-$ray sources in the Galactic plane has been a mystery. The only firmly identified source class has been spin-powered pulsars, whose phase-averaged emission appears constant. However, the $EGRET$ data clearly show that there is at least one class of variable emitters in the Galaxy (eg. Tompkins 1999). Several suggestions as to the nature of these objects have been put forward but no strong observational evidence has previously supported a particular source class. However, the variability timescale of a few months to a few years for sources with emission above 100 MeV gives us some clues to their nature. First, a maximum emission region size can be associated with a light crossing time of a few months or less. For typical Galactic distances of a few kiloparsecs, this implies angular extents on the sky $\\la 1^{\\prime}$. An acceleration mechanism that can generate very high energy particles in a short period must also be present. One known source class fitting these criteria are pulsar wind nebulae (PWN). ", + "conclusions": "" + }, + "0112/astro-ph0112233_arXiv.txt": { + "abstract": "Weak gravitational lensing is a promising tool for the study of the mass distribution in the Universe. Here we report some partial results\\footnote{More complete and detailed results will be published elsewhere} that show how lensing maps can be used to differentiate between cosmological models. We pay special attention to the role of noise and smoothing. As an application, we use mock convergence fields constructed from N-body simulations of the large-scale structure for three historically important cosmological models. Various map analyses are used, including Minkowski functionals, and their ability to differentiate the models is calculated and discussed. ", + "introduction": " ", + "conclusions": "" + }, + "0112/astro-ph0112469_arXiv.txt": { + "abstract": "It is believed that orphan afterglow searches can help to measure the beaming angle in gamma-ray bursts (GRBs). Great expectations have been put on this method. We point out that the method is in fact not as simple as we originally expected. Due to the baryon-rich environment that is common to almost all popular progenitor models, there should be many failed gamma-ray bursts, i.e., fireballs with Lorentz factor much less than 100 --- 1000, but still much larger than unity. In fact, the number of failed gamma-ray bursts may even be much larger than that of successful bursts. Owing to the existence of these failed gamma-ray bursts, there should be many orphan afterglows even if GRBs are due to isotropic fireballs, then the simple discovery of orphan afterglows never means that GRBs be collimated. Unfortunately, to distinguish a failed-GRB orphan and a jetted but off-axis GRB orphan is not an easy task. The major problem is that the trigger time is unknown. Some possible solutions to the problem are suggested. ", + "introduction": "Occurring in the deep Universe, GRBs are the most relativistic phenomena ever known. The standard fireball model (M\\'{e}sz\\'{a}ros \\& Rees 1992; Dermer, B\\\"{o}ttcher \\& Chiang 1999) requires that to successfully produce a GRB, the initial Lorentz factor of the blastwave should typically be $\\gamma_0 \\geq 100$ --- 1000 during the main burst phase (Piran 1999; Lithwick \\& Sari 2000). Generally speaking the requirement of ultra-relativistic motion is to avoid the so called ``compactness problem''. A modest variation in the Lorentz factor will result in a difference of the opacity of the high-energy $\\gamma$-ray photons by a factor of $\\sim 10^3$ (Totani 1999). Additionally, assuming synchrotron radiation, the observed peak frequency is strongly dependent on $\\gamma$, $\\nu_{\\rm m} \\propto \\gamma^4$ (M\\'{e}sz\\'{a}ros, Rees \\& Wijers 1998). Thus a Lorentz factor of $\\gamma_0 \\leq 50$ makes the blastwave very inefficient in emitting $\\gamma$-ray photons. So, to successfully produce a GRB, we need $\\gamma_0 \\geq 100$ --- 1000. However, theoretically it is not easy to construct a model to generate such ultra-relativistic motions. Currently there are mainly two kinds of progenitor models, the collapse of massive stars (with mass $M \\geq 40 M_{\\odot}$), or the collision of two compact stars (such as two neutron stars or a neutron star and a black hole). Since a baryon-rich environment is involved in all these models, some researchers are afraid that the baryon-contamination problem may exist. But this problem maybe is not as serious as we previously expected. Let us first take the collapsar model (MacFadyen \\& Woosley 1999) as an example. We can imagine that the baryon mass and energy released in different collapsar events should vary greatly, then $\\gamma_0$ of the resultant fireballs may also vary in a relatively wide range. In most cases, $\\gamma_0$ should be very low (i.e., $\\gamma_0 \\ll 100$), but there still could be a few cases (e.g., one percent or even one in a thousand) in which the fireball is relatively clean so that the blastwave can be successfully accelerated to $\\gamma_0 \\geq 100$ --- 1000 and produces a GRB. Since the collapsar rate is high enough in a typical galaxy, there should be no problem that such collapsars can meet the requirement of GRB rate (i.e., $10^{-7}$ --- $10^{-6}$ event per typical galaxy per year, under isotropic assumption). Cases are similar in the collisions of two compact stars. In short, we cannot omit an important fact: if GRBs are really due to isotropic fireballs, then there should be much more failed GRBs (i.e., fireballs with Lorentz factor much less than one hundred, but still much greater than unity). These FGRB fireballs can contain similar initial energy as normal GRB fireballs, i.e., $E_0 \\sim 10^{51}$ --- $10^{53}$ ergs, but they are polluted by baryons with mass $M_0 \\sim 10^{-5}$ --- $10^{-3} M_{\\odot}$. Radiation from these FGRBs should mainly be in x-ray bands in the initial bursting phase, not in $\\gamma$-ray bands. In fact, BeppoSAX team has reported the discovery of several anomalous events named as fast X-ray transients, X-ray rich GRBs, or even X-ray-GRBs. They resemble usual GRBs except that they are extremely X-ray rich (Frontera et al. 2000; Kippen et al. 2001; Gandolfi \\& Piro 2001). Observational data on this kind of events are being accumulated rapidly. Recent good examples include GRBs 011030, 011130 and 011211 (Gandolfi et al. 2001a, b; Ricker et al. 2001). We propose that these events are probably just FGRBs. Huang et al. (1998) and Dai, Huang \\& Lu (1999) have pointed out that for afterglow behaviour, the parameter $E_0$ is decisive, while $M_0$ is only of minor importance, especially at late stages. So, FGRBs should also be associated with prominent afterglows. In Figure~1 we compare the theoretical optical afterglows from FGRBs with those from isotropic GRBs and jetted GRBs. We can see that the light curve of an FGRB afterglow differs from that of a successful isotropic burst only slightly, i.e., only notable at early stages. In our calculations, we have used the methods developed by Huang et al.(1999a, b, 2000b, d), i.e., for the dynamical evolution of isotropic fireballs we use \\begin{equation} \\label{dgdmass1} \\frac{{\\rm d} \\gamma}{{\\rm d} m} = - \\frac{\\gamma^2 - 1} {M_0 + \\epsilon m + 2 ( 1 - \\epsilon) \\gamma m}, \\end{equation} where $m$ is the swept-up mass and $\\epsilon$ is the radiation efficience. Eq. (1) has been proved to be proper in both ultra-relativistic phase and non-relativistic phase (Huang, Dai \\& Lu 1999a, b). For jetted ejecta, the following equation is added (Huang et al. 2000b, d), \\begin{equation} \\label{dthetadt2} \\frac{{\\rm d} \\theta}{{\\rm d} t} = \\frac{c_{\\rm s} (\\gamma + \\sqrt{\\gamma^2 - 1})}{R}, \\end{equation} where $R$ is the blastwave radius, and the co-moving sound speed $c_{\\rm s}$ is given realistically by \\begin{equation} \\label{cssquare3} c_{\\rm s}^2 = \\hat{\\gamma} (\\hat{\\gamma} - 1) (\\gamma - 1) \\frac{1}{1 + \\hat{\\gamma}(\\gamma - 1)} c^2, \\end{equation} with $\\hat{\\gamma} = (4 \\gamma + 1) / (3 \\gamma)$ the adiabatic index. In fact, in beamed GRB models, there should also be many FGRBs, i.e., beamed ejecta with $1 \\ll \\gamma_0 \\ll 100$. We call them beamed FGRBs. Afterglow from beamed FGRBs has also been illustrated in Figure~1. In this article emphasises will be put on isotropic FGRBs, so by using ``FGRBs'' we will only mean isotropic FGRBs unless stated explicitly. ", + "conclusions": "To successfully produce a GRB, the blastwave should be ultra-relativistic, with Lorentz factor typically larger than 100 --- 1000. However, in almost all popular progenitor models, the environment is unavoidably baryon-rich. We believe that only in very rare cases can an ultra-relativistic blastwave successfully break out to give birth to a GRB, and there should be much more failed GRBs, i.e., fireballs with Lorentz factor much less than 100 but still much larger than unity. In fact, this possibility has also been mentioned by a number of authors, such as M\\'{e}sz\\'{a}ros \\& Waxman (2001). Owing to the existence of FGRBs, there should be many orphan afterglows even if GRBs are due to isotropic fireballs. Then the simple discovery of orphan afterglows does not necessarily mean that GRBs be highly collimated. To make use of information from orphan afterglow surveys correctly, we should first know how to discriminate a jetted GRB orphan and an FGRB one. This can be done only by checking the detailed afterglow light curve. However, we have shown that the derivation of a satisfactory light curve for an orphan afterglow is difficult. The major problem is that we do not know the trigger time. In Section 3.2, some possible solutions to the problem are suggested. Unfortunately many of these solutions are still quite impractical in the foreseeable future, which means measure of GRB beaming angle using orphan afterglow searches is extremely difficult currently. However, special attention should be paid to the second solution. Usually, FGRBs manifested themselves as fast X-ray transients during the main burst phase, while jetted but off-axis GRBs went unattended completely. If the fast X-ray transients (or X-ray-GRBs) observed by BeppoSAX are really due to FGRBs, then afterglows should be detectable. We propose that this kind of events should be followed rapidly and extensively in all bands, just like what we are doing for GRBs. If observed, afterglows from these anomalous events can be used to check our concept of FGRBs, and even to test the fireball model under quite different conditions (i.e., when $\\gamma_0 \\ll 100$). Also, these FGRBs can provide valuable information for our understanding of GRBs, especially on the progenitor models. Note that beamed FGRBs can also give birth to fast X-ray transients if they are directed toward us, but afterglows from such a beamed FGRB and an isotropic FGRB can be discriminated easily from the light curves (see Figure~1). It is very interesting to note that optical afterglows from two X-ray-GRBs, 011130 and 011211, have been observed (Garnavich, Jha \\& Kirshner 2001; Grav et al. 2001). Their redshifts were measured to be $z = $ 0.5 and 2.14 respectively (Jha et al. 2001; Fruchter et al. 2001), eliminates the possibility that they were ordinary classic GRBs residing at extremely high redshifts ($z \\geq 10$). We propose that they should be FGRBs (either isotropic or beamed) or just jetted GRB ``orphan''. However, the observational data currently available are still quite poor so that we could not determine their nature definitely. As for other X-ray-GRBs without a measured redshift, the possibility that they were at redshifts of $z \\geq 10 $ can not be excluded. Finally, the concept of FGRBs is based on the fact that most popular progenitor models for GRBs are baryon-rich. But cases are quite different for another kind of progenitor models where strange stars are involved. Strange stars, composed mainly of u, d, and s quarks, are compact objects which are quite similar to neutron stars observationally (Alcock, Farhi \\& Olinto 1986). A typical strange star (with mass $\\sim 1.4 M_\\odot$) can have a normal matter crust of less than $\\sim 2 \\times 10^{-5} M_{\\odot}$ (Alcock, Farhi \\& Olinto 1986), or even as small as $\\sim 3 \\times 10^{-6} M_{\\odot}$ (Huang \\& Lu 1997a, b). Then baryon contamination can be directly avoided if GRBs are due to the phase transition of neutron stars to strange stars (Cheng \\& Dai 1996; Dai \\& Lu 1998b) or collisions of binary strange stars. In these models, there should be very few FGRBs." + }, + "0112/astro-ph0112143_arXiv.txt": { + "abstract": "Radio sources from the decametric UTR--2 catalog were cross-identified with other radio catalogs at higher frequencies. We used the CATS database to extract all sources within the UTR beam size ($\\sim$40$'$) to find candidate radio identifications. Using the least squares method we fitted the spectrum of each source with one of a set of curves. We extracted NVSS and FIRST radio images for the radio-identified sources, and looked for a possible relation between size and spectral index. ", + "introduction": " ", + "conclusions": "" + }, + "0112/astro-ph0112375_arXiv.txt": { + "abstract": "It is known from theory that, by means of a plasma physics approach, it is possible to obtain a simple formula to calculate the approximate height of a meteor (Foschini, 1999). This formula can be used in case of forward scatter of radio waves and has the advantage that it does not depend on the diffusion coefficient. On the other hand, it is possible to apply the formula to a particular type of meteor only (overdense meteor type I), which is a small fraction of the total number observed. We have carried out a statistical analysis of several radio echoes from meteor showers recorded during last years by a radio observer located in Belgium. Results are compared and discussed with those obtained with other methods and available in literature. ", + "introduction": "A meteoroid enters the Earth's atmosphere at hypersonic speed and it collides with air molecules. The high kinetic energy involved in the process determine the transformation of a solid body into a plasma, which can scatter radio waves and can emit light (meteor). During sixties and seventies several works investigated the formation and evolution of the meteor, with particular attention to diffusion, in order to study mesospheric winds. A complete review of standard meteor science can be found in Ceplecha et al. (\\cite{CEPLECHA1}). However, there are still some aspects not well understood about the physical properties of a meteor, specifically whether it is an ionized gas or a plasma. During past years, these two appellatives were often used as synonymous in meteor physics, even though they indicate two different states of the matter. In some studies, such as those about diffusion, specific plasma properties are taken into account (e.g. ambipolar diffusion); however in other studies, such as about radiowave scattering, the meteor is simply considered a long narrow column of ionized gas. This can appear as a futile debate, but it hides important concepts. Specifically, a plasma has collective properties (e.g. Langmuir frequency) that an ionized gas has not. A first attempt to study the meteor as a plasma was carried out by Herlofson (\\cite{HERLOFSON}). He investigated the proper oscillations in the meteor and their interaction with radio waves. But, at our knowledge, none continued his studies. Only in 1999 the question of collective oscillations in meteoric plasma was reprised (Foschini \\cite{FOSCHINI1}). Perhaps, this gap may be explained by taking into account that, according to purposes of meteor astronomy, it was sufficient to use the approximation of the long narrow cylinder. However, the meteoric plasma is something more complex than a reflecting rod and it is necessary to study it. There are several types of oscillations and instabilities, which can interact with radio waves. The scattering is not the only process: for example, fluctuations from equilibrium may lead to transformation of waves (longitudinal to transverse and vice versa). The question is: are such processes present in a meteoric plasma? We think that the study of plasma collective oscillations may give new useful tools to understand the physics of meteors. Some basic concepts about meteoric plasma were settled in a previous paper (Foschini 1999), thereafter called Paper I. According to the theory exposed there, radio echoes can be divided into two classes and two subclasses. Then we have underdense and overdense echoes, according to whether the Langmuir frequency is higher or lower than the radio wave frequency. Overdense echoes totally reflect electromagnetic waves, but the presence of binary collisions among ions and electrons weaken the collective oscillations of the plasma, allowing the propagation of the waves, even though with strong attenuation. Therefore, we can divide the overdense echoes into two subclasses: type I, when there is total reflection; type II, when binary collisions allow the propagation. The division between overdense type I and II depends on the electron--ion collision frequency, which in turn depends on electron density and ion cross section. In Paper I, for the sake of simplicity, we considered potassium ion, that is the chemical element with lower ionization energy. In addition, recent studies show that potassium seems to be much more important in the evolution of meteor than previously thought (von Zahn et al. \\cite{MURAD}). With this assuption, the division between overdense echoes occurs at about $10^{17}$~m$^{-3}$. It is worth noting that this border can be moved by considering other elements. But the calculation of particle distribution and evolution in a meteoric plasma will be object of other papers. The overdense type I echoes derives from total reflection of radio waves (see Fig.~\\ref{FIG1} for an example). This allows to calculate the height of the meteor in an easy way, as shown in the Paper I. Here we want to present a statistical sample of several meteor showers, for which we have calculated the height. Data will be discussed and compared with available data in literature. \\begin{figure}[t] \\centering \\includegraphics[scale=0.4]{H2168_f1.eps} \\caption{Example of overdense type I echo.} \\label{FIG1} \\end{figure} ", + "conclusions": "We have carried out an analysis of several overdense radio echoes, recorded during last years by a radio observer located in Belgium. We have analysed a particular class of overdense meteors (type I) and measured height distributions are in good agreement with previous results obtained by McKinley (\\cite{MCKINLEY}), even though only for the secondary peak. We suppose that the first peak in McKinley's work should be due to overdense type II meteors, while the secondary peak, recorded also by our system, appeared to be due to overdense type I meteors. We think that collective properties of the meteoric plasma (Langmuir oscillations) hide some characteristics of the originary cosmic body, specifically there is no clear dependence on speed. Further study, mainly theoretical and able to take into account collective properties of plasma, are required to assess the particle dynamics in the meteor. It should be noted that our studies were carried out with an amateur forward scatter system and we have no full control on it. Moreover, heuristic considerations were introduced in order to minimize uncertainties, but results showed that they were justified by facts. The agreement with previous works, with other techniques, supports our conclusions. The future availability of a full forward scatter system would be of great help in more detailed studies." + }, + "0112/astro-ph0112139_arXiv.txt": { + "abstract": "\\noindent It is shown that the dynamics of cosmologies sourced by a mixture of perfect fluids and self--interacting scalar fields are described by the non--linear, Ermakov--Pinney equation. The general solution of this equation can be expressed in terms of particular solutions to a related, linear differential equation. This characteristic is employed to derive exact cosmologies in the inflationary and quintessential scenarios. The relevance of the Ermakov--Pinney equation to the braneworld scenario is discussed. ", + "introduction": " ", + "conclusions": "" + }, + "0112/astro-ph0112555_arXiv.txt": { + "abstract": "{ We have constructed the foundations to a series of theoretical diagnostic methods to probe the jet phenomenon in young stars as observed at various optical forbidden lines. We calculate and model in a self-consistent manner the physical and radiative processes which arise within an inner disk-wind driven magnetocentrifugally from the circumstellar accretion disk of a young sun-like star. Comparing with real data taken at high angular resolution, our approach will provide the basis of systematic diagnostics for jets and their related young stellar objects, to attest the emission mechanisms of such phenomena. This work can help bring first-principle theoretical predictions to confront actual multi-wavelength observations, and will bridge the link between many very sophiscated numerical simulations and observational data. Analysis methods discussed here are immediately applicable to new high-resolution data obtained with HST and Adaptic Optics.} ", + "introduction": "Herbig-Haro objects (Herbig 1950,Haro 1950) are small nebulous objects that trace jet-like, collimated structures, with characteristic emission spectra of hydrogen and optical forbidden lines of [O~I], [N~II], and [S~II], in the red wavelengths. Mechanisms that produce these phenomena play necessary roles in the making of young stars. Whether jets alone can also drive the often associated molecular outflows has been long debated (Reipurth \\& Bally 2001). The morphology and the distribution of mass and momenta of these outflows argue for wide-angle winds pushing the ejecta and sweeping up materials at wide solid angles (Shu et al 1991). X-winds, by construction, are wide-angle winds, launched magneto-centrifugally near the innermost edge of the circumstellar disks of YSOs (Shu et al 1994), that expand and quickly fill space to ultimately collimate at large distances. Shang et al. (1998) demonstrated the density structure and kinematic information for an x-wind based on a semi-analytic method (Shang 1998), by computing emissions from [S~II]$\\lambda$6716 and [O~I]$\\lambda$6300 lines with uniform ionization conditions throughout the smooth flow without identifying the sources of excitation. The synthetic images for the forbidden lines show strongly cylindrically stratified structures in the center, and suggested the visual appearance of jets is an {\\it optical illusion} out of an intrinsically very divergent flow (Shu et al 1995). It is very compelling to learn whether conditions arising self-consistently in flows driven by magnetically interacting young star-disk systems can reproduce the many observed features known about the dynamics, morphology, and excitation conditions of jets and Herbig-Haro objects. To identify a definitive set of diagnostic tools to help probe the highly collimated jet emissions from the theoretical aspects, is the primary task. This also serves as the first test to bring a well-developed MHD model based on semi-analytic approaches into stringiest confrontation with observations. The fundamental concepts and approaches can be generalized with many multiwavelength observations, and transferred to other theoretical MHD wind/jet models, such as the many interesting numerical simulations discussed in this volume. ", + "conclusions": "In this short article, we summarized recent theoretical developments that lead to direct comparison with real data by first-principle modelling of MHD x-winds and physical processes occurring near the YSOs. We argue that at their base, YSO jets are optical illusions associated with the excitation mechanisms by which atomic forbidden lines are excited. Gentle time-variabilities or pulses implied by knots in observed jets mainly contribute to the variation of thermal excitation on top of background flow, whose densities and velocities are closely maintained by steady state values. This is a demonstration that, first-principle calculations of self-consistent physical processes in young stellar winds provide robust theoretical help on probing the jet phenomena." + }, + "0112/astro-ph0112280_arXiv.txt": { + "abstract": "We analyze 37 {\\it RXTE} observations of the type 2 Seyfert galaxy Mrk348 obtained during a period of 14 months. We confirm the spectral variability previous reported by Smith et al., in the sense that the column density decreases by a factor of $\\sim$ 3 as the count rate increases. Column density variations could possibly originate either due to the random drift of clouds within the absorption screen, or due to photoionization processes. Our modeling of the observed variations implies that the first scenario is more likely. These clouds should lie in a distance of $>$2 light years from the source, having a diameter of a few light days and a density of $>10^7~ \\rm cm^{-3}$, hence probably residing outside the Broad Line Region. ", + "introduction": "Monitoring observations are a powerful tool in the study of the nuclear environment of type 2 Seyfert galaxies. Although the presence of X-ray flux variability in Seyfert-2 galaxies is well established (eg Georgantopoulos \\& Papadakis 2001 and references therein), our understanding of their spectral variations remains limited. Nevertheless, monitoring of the X-ray spectra of a few bright Seyfert-2 galaxies have revealed some intriguing results. For example, Warwick et al. (1988) first reported variations of the column density in ESO 103-G35. Using {\\it EXOSAT} data, they found a decrease in the column density by a factor of $\\sim1.7$ in a period of 90 days. Warwick et al. (1993) found column density variations in NGC 7582 by a factor of $\\sim$ 3 over an interval of about 4 years, using \\ginga data. They attributed these variations to motions of clouds near the central source. Investigation of the spectrum of NGC 7582 based on {\\it ASCA} data by Xue et al. (1998) confirmed the existence of significant column density variations (by a factor of $\\sim 2$) over a timescale of 2 years. Variability analysis of column density variations in large samples of Seyfert-2 galaxies has been performed by Malizia et al. (1997) and Risaliti, Elvis \\& Nicastro (2002) using mainly literature data. They find variation of the column density in time scales of a few months up to several years. Systematic monitoring observations only became feasible with the {\\it RXTE } mission. In particular, Georgantopoulos et al. (1999), Georgantopoulos \\& Papadakis (2001), Smith, Georgantopoulos \\& Warwick (2001) present monitoring observations of several Seyfert-2 galaxies (Mrk3, ESO 103-G35, IC 5063, NGC 4507, NGC 7172 and Mrk 348) spanning time periods from about seven days to seven months. They found statistically significant spectral variations in all cases. In some objects the variations appear to be caused by intrinsic power-law slope changes whereas in others column density variations dominate. In this paper we present an analysis of 37 {\\it RXTE} observations of the Seyfert 2 galaxy Mrk348. Previously Smith et al. (2001) have reported the results from 12 $RXTE$ observations covering a period of six months. Here we use an expanded sample of 37 $RXTE$ monitoring observations of Mrk 348 (including the 12 reported by Smith et al. 2001), spanning a time interval of 14 months, to investigate further the nature of the X-ray spectral variability exhibited by this source. More specifically, we investigate whether photoionization of the absorbing screen or alternatively motion of clouds along the line of sight can reproduce the observed spectral variability. ", + "conclusions": "We model the X-ray spectral variability of Mrk348 by analyzing 37 RXTE observations over a period of 14 months. The present work extends our previous study (Smith et al. 2001) where 12 observations spanning a period of 6 months were used. We find that the column density decreases by a factor of 3 with increasing flux, confirming our previous results. These column density variations could arise either due to the random drift of clouds within the absorption screen or due to photoionisation processes. Our modelling shows that photoionization alone cannot easily reproduce the observed column density variations. In particular, both the {\\sl ABSORI} and the {\\sl XSTAR} models yield a worse fit to the data compared to the double neutral absorber model. This is further supported by the {\\it ASCA} observation which does not show any significant evidence for an ionized edge. Alternatively a model which assumes that a neutral or a weakly ionized cloud is moving in front of the source can reproduce successfully the column density variations. If we assume that the observed $N_H$ variations are caused by a single spherical cloud which moves in a circular orbit, then, this cloud should lie in a distance of $>$2 light years from the source having a density of $>1 \\times 10^7 \\rm cm^{-3}$. These values are probably characteristic of a region outside the BLR. Future observations with the XMM mission, which combines large effective area and excellent spectral resolution will shed more light on the origin of the obscuring screen in Seyfert 2 galaxies." + }, + "0112/hep-ph0112060_arXiv.txt": { + "abstract": "\\noindent What can we learn from solar neutrino observations? Is there any solution to the solar neutrino anomaly which is favored by the present experimental panorama? After SNO results, is it possible to affirm that neutrinos have mass? In order to answer such questions we analyze the current available data from the solar neutrino experiments, including the recent SNO result, in view of many acceptable solutions to the solar neutrino problem based on different conversion mechanisms, for the first time, using the same statistical procedure. This allows us to do a direct comparison of the goodness of the fit among different solutions, from which we can discuss and conclude on the current status of each proposed dynamical mechanism. These solutions are based on different assumptions: (a) neutrino mass and mixing, (b) non-vanishing neutrino magnetic moment, (c) the existence of non-standard flavor-changing and non-universal neutrino interactions and (d) the tiny violation of the equivalence principle. We investigate the quality of the fit provided by each one of these solutions not only to the total rate measured by all the solar neutrino experiments but also to the recoil electron energy spectrum measured at different zenith angles by the Super-Kamiokande collaboration. We conclude that several non-standard neutrino flavor conversion mechanisms provide a very good fit to the experimental data which is comparable with (or even slightly better than) the most famous solution to the solar neutrino anomaly based on the neutrino oscillation induced by mass. ", + "introduction": "\\label{sec:intro} Solar neutrino observations coming from Homestake~\\cite{Lande:nv}, Kamiokande~\\cite{Fukuda:1996sz}, SAGE~\\cite{Abdurashitov:uu}, GALLEX~\\cite{Hampel:1998xg}, GNO~\\cite{Altmann:2000ft}, and Super-Kamiokande (SK)~\\cite{Fukuda:1998rq} have been suggesting a picture which is conflicting with the predictions from the Standard Solar Model (SSM)~\\cite{JNB,BP00,concha-novo}, strongly indicating disappearance of solar electron neutrinos on their way from the Sun to the terrestrial detectors. This has been known for many years as the {\\em Solar Neutrino Problem} (SNP)~\\cite{JNB}. The extraordinary new result from the Sudbury Neutrino Observatory (SNO)~\\cite{SNO}, inaugurates a new era in the quest for the solution to the long-standing puzzle of missing solar neutrinos. For the first time in the history of solar neutrino observations, a direct indication of the presence of non-electron active neutrino component in the solar neutrino flux is obtained. This cannot be explained by any conceivable modification of the SSM but do require some departure from the standard electroweak theory. The indication of non-electron active component is based on the difference of $^8$B neutrino flux detected through charged current events in SNO and the neutrino electron elastic scattering events observed by the SK collaboration, the former obtaining a lower rate than the latter. Such difference can be explained by the conversion of electron neutrinos into active non-electron ($\\nu_\\mu$ or $\\nu_\\tau$) neutrinos along their trajectory from the Sun to the detectors at the Earth\\cite{SNO,fogli_sno}. As we will see, the hypothesis of no flavor conversion of solar electron neutrinos is strongly in conflict with the prediction of the SSM; it is now only acceptable at very small confidence level ($\\sim 10^{-12}$) if only the total rates from solar neutrino experiments is considered (see Sec.\\ \\ref{sec:resul}). The disagreement among the observed solar neutrino data and the theoretical predictions can be relaxed to 4~$\\sigma$ level ($7 \\times 10^{-5}$)~\\cite{Bahcall2001} if one allows all the solar neutrino fluxes to be free parameters in fitting the measured solar neutrino event rates. However, this can only be obtained under the extreme assumption of vanishing $^{7}$Be neutrino flux, which is quite difficult to explain. Several mechanisms can induce neutrino flavor conversion when one assumes that neutrinos are endowed with some properties not present in the minimal standard electroweak theory~\\cite{30years}. The most well known mechanism is the neutrino oscillation induced by mass and mixing~\\cite{MNS,vacuum,wolf,MS}. The fact that the terrestrial experiments are less sensitive to these resulting non-electron neutrinos can explain their observed lower counting rates. The purpose of this article is to compare quantitatively the capabilities of several different mechanisms to explain solar neutrino data, for the first time, based on the same statistical procedure. Readers are invited to take a look at Tables \\ref{tab:rates} and \\ref{tab:combined} which summarize our most important results. A number of possible solutions to the solar neutrino anomaly still survives even after SNO results, fitting the data with significant confidence level. In fact, combined analysis of the data, which includes not only the total rates measured by all solar neutrino experiments but also some information which is independent of the total neutrino flux, namely, the energy spectrum and the zenithal dependence of the data, suggests that the large mixing angle MSW solution in matter, as well as the mechanisms based on resonant spin-flavor precession, non-standard neutrino interactions and violation of the equivalence principle all provide a fit of the data with the confidence level $\\gsim$ 60\\%. In Sec.\\ \\ref{sec:mec}, we briefly review several mechanisms that induce flavor conversion of solar neutrinos, which will be discussed in this work: (a) mass induced oscillation in vacuum and in matter (b) resonant spin-flavor precessions induced by a non-vanishing neutrino magnetic moment, (c) the existence of non-standard neutrino interactions inducing flavor-changing and non-universal currents and (d) the violation of the equivalence principle. In Sec.\\ \\ref{sec:proc}, the procedure of the statistical analysis used in this work is presented, while our results are given in Sec.\\ \\ref{sec:resul}. Finally, Sec.\\ \\ref{sec:conc} is devoted to the discussion of our results. ", + "conclusions": "\\label{sec:conc} Even though the conventional mass induced oscillation mechanism, which is theoretically well motivated, can be considered as the most plausible solution to the solar neutrino problem, it is important to realize that solutions based on New Physics in the neutrino sector, such as large neutrino magnetic moment, neutrino flavor-changing and non-universal processes, violation of the equivalence principle, can still be viable ones providing a fit to the solar neutrino data which is comparable to the one based on the conventional mass-induced neutrino oscillation, as we showed in this work. We also note that some of these mechanisms which do not require neutrino mass in the fit to the solar neutrino data do have some close relation with neutrino mass generation. For example, in our phenomenological approach, we ignored neutrino mass in the solution based on non-standard neutrino interactions but there is no available model that prevents neutrinos to acquire mass at radiative level, if flavor changing and flavor non-universal interactions with quarks are present, even if no tree-level mass terms appear in the model. It is fundamental to mention that it is difficult to explain the atmospheric neutrino problem~\\cite{AN} as well as the LSND anomalies~\\cite{LSND} by these alternative mechanisms~\\cite{failed}. Furthermore, while mass induced oscillation and resonant spin-flavor precession solutions to the SNP can be easily conciliated with the standard neutrino oscillation solution to the atmospheric neutrino problem, it is not a trivial task to answer if the non-standard flavor-changing and non-universal neutrino interactions and violation of equivalence principle solutions to the SNP modify or even damage this standard solution to the atmospheric neutrino anomaly~\\cite{hybrid}. Having this picture in mind we conclude that the no specific solution is preferred by the current solar neutrino data, although some solutions may have difficulties in reconciling atmospheric neutrino observations. We emphasize that the solar neutrino observations alone can not yet conclude if neutrinos have non-vanishing mass or magnetic moment. Furthermore, no stringent limit on the existence of NSNI nor on VEP can be currently set by the solar neutrino data. The ultimate goal of the solar neutrino observations is, of course, to perform a direct experimental identification of the solution. For this purpose, we will have to wait for the upcoming solar neutrino experiments, which hopefully will provide a lot of new informations, to reveal the true nature of the flavor conversion mechanism which is behind the SNP. For instance, low energy solar neutrino experiments, such as Borexino~\\cite{borexino}, can possibly discriminate among the solutions considered in this work, as can be seen in Fig. 7 of Ref.~\\cite{nuno-lownu}. In the near future, a new reactor experiment, KamLAND~\\cite{kamland}, may measure $\\nu_e$-disappearance. For all mechanisms studied here, with exception of LMA, we expect no significant disappearance in KamLAND because the baseline is too short to developed oscillation. Then a positive evidence for neutrino oscillation in KamLAND can establish LMA solution. For the negative evidence no conclusion can be drawn to favor a particular mechanism studied in this work. A discussion about the consequences of KamLAND experiment on the mechanism of neutrino mass and mixing hypothesis can be found in Ref.~\\cite{gouvea}. \\vskip 0.5cm Summarizing our conclusions, \\begin{enumerate} \\item The present solar neutrino data only by themselves can not discard, {\\em a priori}, any of the solutions discussed here: (a) neutrino mass and mixing, (b) non-vanishing neutrino magnetic moment, (c) the existence of non-standard flavor-changing and non-universal neutrino interactions and (d) the tiny violation of the equivalence principle. We refer to Tables \\ref{tab:rates} and \\ref{tab:combined} for comparison. All solutions have confidence level over $60\\%$ C.L. providing a very good fit to the solar neutrino data; \\item A very robust statement is that LMA is the best preferred solution for mass induced oscillation scenario, whereas the SMA is the worst one; \\item Future experiments could test these different scenarios and possibly discard some of them. \\end{enumerate}" + }, + "0112/astro-ph0112212_arXiv.txt": { + "abstract": "{ We present new data of the dwarf galaxy \\object{NGC~1569} at 450 $\\mu$m, 850 $\\mu$m and 1200$\\mu$m taken with SCUBA at the JCMT and the bolometer array at the IRAM 30m telescope. After including data from IRAS at 12, 25, 60 and 100 $\\mu$m, we have successfully fitted the dust grain population model of D\\'esert et al. (1990) to the observed midinfrared-to-millimeter spectrum. The fit requires a combination of both large and very small grains exposed to a strong radiation field as well as an enhancement of the number of very small grains relative to the number of large grains. We interpret this as the consequence of large grain destruction due to shocks in the turbulent interstellar medium of \\object{NGC~1569}. The contribution of polyaromatic hydrocarbons (PAH's) is found to be negligible. Comparison of the dust emission maps with an HI map of similar resolution shows that both dust and molecular gas distributions peak close to the radio continuum maximum and at a minimum in the HI distribution. From a comparison of these three maps and assuming that the gas-to-dust mass ratio is the same everywhere, we estimate the ratio of molecular hydrogen column density to integrated CO intensity to be about 25 -- 30 times the local Galactic value. The gas-to-dust ratio is 1500 -- 2900, about an order of magnitude higher than in the Solar Neighbourhood. ", + "introduction": "Dwarf galaxies characteristically have low metallicities and consequently low dust abundances. In dwarf galaxies, both dust properties and amounts differ from those in spiral galaxies, as for instance implied by a difference in IRAS colours (cf. Melisse \\& Israel \\cite{melisse}). However, actual dust abundances and dust composition are only poorly known for dwarf galaxies. The dust mass of a galaxy can be estimated reliably only if good (sub)mm measurements allow to constrain the amounts of relatively cold dust which may dominate the total dust mass with only a very limited contributions to the emission at infrared wavelengths. Such data are available for a limited number of galaxies, and are especially scarce for faint dwarf galaxies. \\object{NGC~1569} (Arp~210; VII~Zw~16) is a nearby irregular dwarf galaxy at a distance of 2.2 Mpc (Israel 1988, hereafter \\cite{israel88}). It is a member of the low-galactic latitude IC~342/Maffei~1/Maffei~2/Dw~1 group, containing at least 15 dwarf galaxies (Huchtmeier et al. \\cite{huchtmeier}). \\object{NGC~1569} is presently in the aftermath of a massive burst of star formation (\\cite{israel88}; Israel \\& De Bruyn \\cite{israel-bruyn}; Waller \\cite{waller}). { Its present star formation rate, derived from the H$\\alpha$ luminosity, is 0.4 \\msun yr$^{-1}$ (Waller \\cite{waller}). In the recent past, this galaxy has experienced a starburst which started about $1 - 2 \\times 10^7$ yr ago as estimated by Israel (\\cite{israel88}) and about $1.5 \\times 10^7$ yr ago determined by Vallenari \\& Bomans (\\cite{vallenari}) from colour-magnitude diagrams. The end of the starburst, about 5 Myr ago, can be well dated by a kink in the synchrotron spectrum (Israel \\& De Bruyn \\cite{israel-bruyn}) and photometric studies (Vallenari \\& Bomans \\cite{vallenari}, Greggio et al \\cite{greggio}). } Although a small galaxy with neutral atomic hydrogen (HI) dimensions of 3 x 2 kpc (Israel \\& Van Driel \\cite{israel90}; Stil \\& Israel, in preparation, hereafter \\cite{stil01}), \\object{NGC~1569} contains two extremely compact luminous star clusters A and B (Ables \\cite{ables}; Arp \\& Sandage \\cite{arp}; Aloisi et al. \\cite{aloisi} and references therein) with bolometric luminosities of order 10$^{8}$ L$_{\\odot}$ located in a deep HI minimum. A third such cluster is embedded in bright emission nebulosity (F.P. Israel $\\&$ W. Wamsteker, unpublished; Prada et al. \\cite{prada}; ; listed as No. 10 by Hunter et al. \\cite{hunter}) coincident with the peak of the radio continuum distribution. The HI observations show the galaxy to be in solid-body rotation out to a radius of 1$'$ (0.64 kpc); however, much of the HI is in chaotic motion (\\cite{stil01}). To the east of \\object{NGC~1569}, an apparently counterrotating HI cloud connected by an HI bridge to \\object{NGC~1569} is observed (Stil $\\&$ Israel \\cite{stil98}). The brightest HI with column densities N(HI) $\\approx 2 \\times 10^{20}$ cm$^{-2}$, occurs in the three main peaks of an HI ridge in the northern half of the galaxy. About a third of the total mass of \\object{NGC~1569} resides in its neutral atomic hydrogen (\\cite{israel88}, \\cite{stil01}). Weak CO emission is found close to the radio continuum, between two HI maxima. It was mapped in the $J$=1--0 and $J$=2--1 transitions by Greve et al. (\\cite{greve}). Comparison with the $J$=1--0 CO detection in a larger beam by Taylor et al. (\\cite{taylor98}) suggests that the maps by Greve et al. contain virtually all CO emission from the western half of \\object{NGC~1569}. Aperture synthesis maps of the same two CO transitions with the IRAM interferometer show distinct CO clouds of sizes ranging from 40 to 100 pc (Taylor et al. \\cite{taylor99}). As these maps recover only a quarter of the single-dish flux, weak and more diffusely distributed CO must be present. CO(3-2) observations show a high $J$=3--2/$J$=2--1 ratio of 1.4 (M\\\"uhle et al. \\cite{muehle}) indicating a warm molecular gas phase. The far-infrared emission from \\object{NGC~1569}, observed with IRAS, is remarkably strong for a dwarf galaxy; the continuum spectrum indicates that the dust in the galaxy is exposed to intense radiation fields (\\cite{israel88}). In this paper we extend the far-infrared spectrum of \\object{NGC~1569} to (sub)millimeter wavelengths by presenting new observations obtained with the IRAM and JCMT telescopes. These new observations allow us, for the first time, to determine the amount of dust in this galaxy and to study its properties in detail. ", + "conclusions": "\\begin{enumerate} \\item We present new maps of the dwarf galaxy \\object{NGC~1569} at 450, 850 and 1200 $\\mu$m taken with SCUBA at the JCMT and the MPIfR bolometer array at the IRAM 30m telescope. Integrated flux-densities at these wavelengths may be compared to those at at 12, 25, 60 and 100 $\\mu$m obtained earlier with IRAS. \\item The steep rise in intensity from 12 to 25 $\\mu$m excludes a significant contribution from PAH's, as the broadband spectrum of these decreases in wavelength. \\item The (sub)millimetre flux densities are high compared to the flux densities at shorter wavelengths and the (sub)millimetre spectrum has a relatively shallow slope ($\\simeq \\lambda^{-2.5}$). Such a spectral shape can be explained by the presence of a significant amount of cold dust. Fits to the observed spectrum by three-temperature dust models, however, require most dust to be at temperature of only about 7 K. \\item We do not favour this explanation. The intense radiation field and low metallicity of \\object{NGC~1569}, implying poor shielding, render it very unlikely that large amounts of dust at such low temperatures may exist in \\object{NGC~1569}. { We show that the high dust opacities necessary to shield a large fraction of the dust from this radiation field are not present.} Even if it could, the resulting gas-to-dust mass ratio would have to be about 100, again an unusually low value for a low metallicity galaxy such as \\object{NGC~1569} ($12+\\log(O/H)=8.19$). \\item Alternatively, the (sub)millimetre flux densities may be dominated by emission of VSG's at various non-equilibrium temperatures. The combined emission is characterized by a wavelength dependence mimicking an extinction coefficient $k_\\lambda \\propto \\lambda^{-1}$. Using the dust model of \\cite{desert}, we achieve good fits for dust exposed to radiation fields similar in spectral shape to the Solar Neighbourhood field but with sixty times higher intensity, as is appropriate for \\object{NGC~1569} (\\cite{israel88}). \\item The fits require VSG to large grain abundances to be enhanced by a factor of 7 as compared to those in Solar Neighbourhood. The precise value of the enhancement factor is slightly model dependent; use of different radiation field yielding less good, but still acceptable fit results, yields slightly different factors. A robust conclusion is that the VSG enhancement factor in \\object{NGC~1569} is 2 -- 7. \\item Although in these models much of the emission originates from very small grains, virtually all of the mass still resides in the large grains. The gas-to-dust mass ratio is 1500 -- 2900, about an order of magnitude higher than in the solar neighbourhood. \\item Both the dust and molecular gas distributions peak at a local minimum in the HI ridge. If gas-to-dust mass ratio are constant over \\object{NGC~1569}, the lack of HI at the very peak of the dust emission indicates the presence of a significant column of molecular gas. From the inferred molecular gas column density and the observed CO emission, we estimate the hydrogen column density to integrated CO intensity conversion factor $X \\approx 5 \\times 10^{21} $ H$_{2}$ mol cm$^{-2}$ (K km s$^{-1})^{-1}$, or about 25 -- 30 times the local Galactic value. \\end{enumerate}" + }, + "0112/astro-ph0112024_arXiv.txt": { + "abstract": "We present hydrodynamical simulations, using a 2-D two fluid model, of bow shocks in a representative regime for pulsar wind driven bow-shock nebulae. We also investigate the behaviour of a passive toroidal magnetic field wounded around the pulsar velocity direction. Moreover we estimate the opacity of the bow-shock to penetration of ISM neutral hydrogen. Finally we compare these numerical results with those from an analytical model. ", + "introduction": "Pulsars moving supersonically with respect to the ambient medium are expected to give rise to a bow shock. In the case of interaction with the interstellar medium (ISM) the \\halpha emission from the nebula may be detected. Only four bow shocks have been discovered so far: PSR~1957+20 (Kulkarni \\& Hester 1988), PSR~2224+65 (Cordes, Romani \\& Lundgren. 1993), PSR~J0437+4715 (Bell 1995), and PSR~0740--28 (Jones, Stapper \\& Gaensler 2001). These nebulae may be detected in optical Balmer lines (mostly \\halpha) as a signature of a non-radiative shock moving through a partially neutral medium (Chevalier \\& Raymond 1980). A puzzling point is that two out of the four known nebulae show a peculiar shape with a conical tail. Various hypothesis have been put forward to justify the conical tail shape: a peculiar ISM density distribution in the surroundings of the PSR~2224+65; effects like mass loading due to neutral atoms which might penetrate the external layers of shocked material (Bucciantini \\& Bandiera 2001, hereafter Paper I) . The nebula near PSR~0740--28, seems to represent an intermediate case, with a ``standard'' head and a conical tail. The ISM is seen by the pulsar as a plane-parallel flow, likely with a constant density, and lasting long enough to produce a steady-state regime. Moreover the ISM is typically partly neutral and the neutral atoms may have collisional mean free paths comparable with the scale length of the system, then invalidating a purely fluid treatment. However as demonstrated in a previous paper (Paper I), if we suppose that the H atoms can be ionized only via collisions, for a large amount of pulsars the presence of a neutral component in the ISM can be neglected . The interaction of the relativistic magnetized pulsar wind with the ISM ionized component is due to the pulsar wind magnetic field compressed on the head of the nebula. The mean free path of particles is typically much smaller than the typical dimension of the nebula: $d_{o}=\\sqrt{{\\cal L}/4\\pi c \\rho_{o} V_{o}^{2}}$,where ${\\cal L}$ and $V_{o}$ are the pulsar luminosity and velocity, and $\\rho_{o}$ the density of the dragged component of ambient medium (Paper I). ", + "conclusions": "" + }, + "0112/astro-ph0112354_arXiv.txt": { + "abstract": "We review recent theoretical developments on pulsar winds, their nebulae and relativistic shock acceleration, and show how they illuminate unsolved problems in plerion spectra, in particular the multiple spectral breaks in the Crab and the low-frequency breaks of plerions such as G\\,21.5--0.9 and 3C\\,58. Recent work on Fermi acceleration theory at relativistic shocks shows that a particle spectral index of 2.2--2.3, compatible with the X-ray spectra of plerions, results under a wide variety of assumptions. If pulsar winds contain ions as well as electrons and positrons, the mechanism of Hoshino et al.\\ (1992), which yields harder spectra, would operate at lower energies and may explain the flat radio spectral indices of plerions. This scenario implies wind parameters in the Crab compatible with the pulsar wind acceleration model of Lyubarsky \\& Kirk (2001). Recent hydrodynamical simulations of plerion evolution inside SNR blast waves demonstrate that the passage of the reverse shock rapidly compresses the plerion. Using a simple isobaric model, we investigate the influence of the resulting magnetic field compression and decrease in shock radius on the evolution of the plerion spectrum. We suggest that the passage of the reverse shock may explain the low-frequency breaks in 3C\\,58 and G\\,21.5--0.9, as well as the increase in 3C\\,58's radio flux. ", + "introduction": "\\subsection{Observations and Synchrotron Cooling} Plerions are characterized in X-rays by hard, nonthermal power-law spectra. In the case of the Crab Nebula, where statistics are best, the total integrated spectrum has a best-fit power-law index $\\alpha_X = 1.1$ (Toor \\& Seward 1974) in energy ($F_\\nu \\propto \\nu^{-\\alpha}$, corresponding to a photon index $\\Gamma_X \\equiv \\alpha_X + 1 = 2.1$). Within their larger uncertainties, the X-ray power-law indices of most other plerions appear compatible with this value. The recent availability of spatially resolved spectra of plerions with {\\em Chandra} and {\\em XMM} reveals spectral steepening towards the edges, with the hardest spectrum at the center having an index around $\\Gamma_X \\approx 1.6$. The steepening is indicative of synchrotron cooling of a centrally injected hard power-law distribution of electrons and positrons, as is the difference of 0.5 between the central and spatially integrated spectral indices. The corresponding injected particle spectral index is $p = 2.2$, defined by \\beq \\dot{N}(\\gamma) \\,{\\rm d} \\gamma \\propto \\gamma^{-p} \\,{\\rm d} \\gamma \\eeq where $\\dot{N}$ is the injection rate and $\\gamma$ the particle Lorentz factor. \\subsection{Theory: Fermi Acceleration at Ultra-Relativistic Shocks} This non-thermal population of electrons and positrons is generally assumed to be accelerated at the termination shock of a highly relativistic pulsar wind. Recent investigations of Fermi acceleration at such relativistic shocks (Bednarz \\& Ostrowski 1998; Gallant \\& Achterberg 1999; Kirk et al.\\ 2000; Achterberg et al.\\ 2001) have shown that the resulting spectra, in the limit of high Lorentz factors and of a turbulent magnetic field downstream, have power-law indices $p$ in the range 2.2--2.3 for a variety of transport assumptions, compatible with the above inferred value for the injected spectrum. While further investigations, in particular using 3-D plasma simulations, are needed to confirm that the required levels of magnetic turbulence can be achieved, it seems reasonable to identify the acceleration mechanism for X-ray emitting electrons with Fermi acceleration at the pulsar wind termination shock. Further observational evidence for this scenario comes from gamma-ray burst afterglows, whose spectra can also be explained in terms of Fermi acceleration at the highly relativistic outer blast wave (e.g. Gallant et al.\\ 2000). ", + "conclusions": "The X-ray spectra of plerions are compatible with Fermi acceleration at ultra-relativistic shocks, which yields a power-law distribution of injected particles with spectral index $p_X = 2.2$--2.3, above a critical energy $\\gamma_{\\rm crit}$. Plerion radio spectra seem compatible with resonant ion cyclotron wave acceleration, yielding a harder power-law index, $p_r = 1$--1.6, and fixing the break energy at $\\gamma_{\\rm crit} \\sim (m_i/m_e) \\gamma_{\\rm sh}$. In the case of the Crab Nebula, identification of the synchrotron cooling break in the FIR and the injection break in the UV inplies a wind Lorentz factor of about $10^3$, in sharp contrast with the model of Kennel \\& Coroniti (1984). This value of the wind Lorentz factor is compatible with the striped wind model of Lyubarsky \\& Kirk (2001). Finally, compression by the reverse shock might be responsible for the low-frequency breaks observed in G\\,21.5--0.9 and 3C\\,58, among others, but this requires a dense surrounding medium. One prediction of our scenario for particle acceleration is that small-scale features near the wind termination shock, where synchrotron losses have not had time to operate, should reflect the unsteepened injection spectrum, with a single spectral break. The Crab Nebula's wisps should be just such features, and comparison of the frequency of this injection break in the wisps and the Nebula as a whole would then allow a determination of the relative magnetic field values." + }, + "0112/astro-ph0112481_arXiv.txt": { + "abstract": "{I found one long X-ray flare from the X-ray burster GX\\,3+1 in almost 6~years of observations with the RXTE All Sky Monitor (ASM). The event had a peak flux of about 1.1~Crab (1.5-12\\,keV), lasted between 4.4 and 16.2~hours and exhibited a fluence of more than about 5$\\times$10$^{41}$\\,erg for a source distance of 5\\,kpc. During the exponential-like decay, with an exponential decay time of 1.6~hours, spectral softening is seen. The total ASM effective exposure time on GX\\,3+1 is estimated to be around a year. The flare bears all the characteristics of the recently discovered so-called superbursts in other X-ray burst sources. ", + "introduction": "Recently, six long X-ray flares lasting several hours have been identified in five low-mass X-ray binaries (Cornelisse et al.\\ 2000, 2001, Strohmayer \\&\\ Brown 2001; Wijnands 2001; Kuulkers et al.\\ 2001). For one source, 4U\\,1636$-$53, two such events were reported which occurred about 4.7~years after each other (Wijnands 2001). So far, these flares have only been seen in X-ray bursters with persistent pre-flare luminosities of $\\simeq$0.1--0.3 times the Eddington luminosity (Wijnands 2001; Kuulkers et al.\\ 2001). The long X-ray flares share many of the characteristics of type~I X-ray bursts\\footnote{Type~I X-ray bursts have light curves with a rise which is faster than the exponential-like decay; their emission is well described by black-body radiation with temperatures, $kT$, around 2\\,keV and apparent black-body radii around 10\\,km; they show X-ray spectral softening during the decay. They have durations of seconds to minutes. For a review, see Lewin et al.\\ (1993).} and are, therefore, attributed to thermonuclear runaway events on a neutron star (e.g.\\ Cornelisse et al.\\ 2000). The differences with type~I bursts are their long duration (exponential decay times of a few hours), their large fluences (about 10$^{42}$\\,erg), and their rarity. Because of the large fluences, the flares are referred to as `superbursts'. A likely fuel for the superbursts is carbon, left over from stable and unstable hydrogen and/or helium burning (Cumming \\&\\ Bildsten 2001; Strohmayer \\&\\ Brown 2001). Unstable electron capture by protons with subsequent capture of the resulting neutrons by heavy nuclei is another conceivable option (Kuulkers et al.\\ 2001). I here report on a superburst seen with the {\\it Rossi X-ray Timing Explorer} (RXTE) All Sky Monitor (ASM) from GX\\,3+1 in June 1998. For a preliminary announcement of this event see Kuulkers (2001). The overall X-ray intensity of the GX 3+1 varies slowly on time scales of months to years by a factor of about 2 (e.g.\\ Makishima et al.\\ 1983; see Fig.~\\ref{plot_asm}a). Type~I X-ray bursts in GX\\,3+1 were first discovered by Hakucho (Makishima et al.\\ 1983). A type~I X-ray burst with radius expansion (due to the burst luminosity reaching the Eddington limit) was observed by the Proportional Counter Array (PCA) onboard RXTE, enabling one to estimate the distance to the source to be in the range 4--6\\,kpc (Kuulkers \\&\\ van der Klis 2000). ", + "conclusions": "I found an X-ray flare from GX\\,3+1 in RXTE/ASM data spanning about 6~yrs; the flare had a decay time of 1.6\\,hr and a duration of longer than 4.4\\,hr, but shorter than 16.2~hr. During the exponential-like decay the flare spectrum softened. I conclude that the flare has its origin in unstable thermonuclear burning. The total fluence of the event is between about 5$\\times$10$^{41}$\\,erg and 2$\\times$10$^{42}$\\,erg for a distance of 5\\,kpc. The maximum net-burst flux reached during a normal type~I X-ray burst seen by Kuulkers \\&\\ van der Klis (2000) was 6.9$\\times$10$^{-8}$\\,erg\\,cm$^{-2}$\\,s$^{-1}$. Since this burst was a radius-expansion event the corresponding luminosity at the neutron star surface reached the Eddington limit. The flare, therefore, had an observed maximum of about 0.5 times the Eddington value (but note that the actual peak of the flare was missed). The persistent ASM count rate before and after the flare was similar to that observed around the radius expansion burst, indicating the persistent flux near the X-ray flare was about 0.2 times the Eddington value. The long X-ray flare bears all the characteristics of the superbursts discovered recently in five other X-ray burst sources: 4U\\,1735$-$44, Ser\\,X-1, 4U\\,1636$-$53, 4U\\,1820$-$30 and KS\\,1731$-$260 (Cornelisse et al.\\ 2000, 2001; Wijnands 2001; Strohmayer \\&\\ Brown 2001; Kuulkers et al.\\ 2001). They have been shown to be thermonuclear events on the surfaces of neutron stars, in which carbon (Cumming \\&\\ Bildsten 2001; Strohmayer \\&\\ Brown 2001) or electron capture (Kuulkers et al.\\ 2001) plays the dominant role. Assuming the persistent luminosity is a direct measure of the mass accretion rate, $\\dot{M}$, onto the neutron star, one would infer $\\dot{M}$$\\simeq$2$\\times$10$^{17}$\\,g\\,s$^{-1}$ in GX\\,3+1 just before the superburst. But since the persistent luminosity is variable by a factor of $\\simeq$2 and at the time of the superburst it was at a minimum as seen by the ASM in $\\simeq$6~yrs, I assume an average $\\dot{M}$ of $\\simeq$3$\\times$10$^{17}$\\,g\\,s$^{-1}$. If the ignition of the superburst occurs in a carbon-rich layer, recurrence times of $\\simeq$13~yrs are then expected, unless the ignition is prematurely triggered (Strohmayer \\&\\ Brown 2001). Note that this case is applicable to a neutron star accreting pure helium, such as 4U\\,1820$-$30. Of the other systems, only 4U\\,1735$-$44 and 4U\\,1636$-$53 are known to have hydrogen-rich donors (Augusteijn et al.\\ 1998); this is not as yet clear for GX\\,3+1. For hydrogen/helium accreting neutron stars it has been shown that carbon may ignite earlier when in an ocean of heavy elements, leading to a smaller recurrence time of $\\simeq$2~yrs at the average $\\dot{M}$ (Cumming \\&\\ Bildsten 2001). If, on the other hand, the superbursts are due to unstable electron capture by protons with subsequent capture of the resulting neutrons by heavy nuclei, then recurrence times of less than 0.5~yr are expected (Kuulkers et al.\\ 2001). The expected recurrence times in the latter two cases are more or less compatible with seeing one superburst from GX\\,3+1 in a total effective exposure time of about 1 to 1.5~yrs, as well as seeing two such events 4.7~yrs apart from 4U\\,1636$-$53 (Wijnands 2001). Effective exposure times can in principle be determined for the other superburst sources, using RXTE/ASM data together with data from other instruments (e.g.\\ BeppoSAX/WFCs); however, this is outside the scope of this paper. Longer and more frequent monitoring of these sources will enable one to discriminate the origin of the superbursts." + }, + "0112/astro-ph0112162_arXiv.txt": { + "abstract": "We compare the amplitudes of fluctuations probed by the 2dF Galaxy Redshift Survey and by the latest measurements of the Cosmic Microwave Background anisotropies. By combining the 2dFGRS and CMB data we find the linear-theory rms mass fluctuations in $8 \\Mpc$ spheres to be $\\sigma_{8{\\rm m}} = 0.73 \\pm 0.05$ (after marginalization over the matter density parameter $\\omegam$ and three other free parameters). This normalization is lower than the COBE normalization and previous estimates from cluster abundance, but it is in agreement with some revised cluster abundance determinations. We also estimate the scale-independent bias parameter of present-epoch $L_s = 1.9L_*$ APM-selected galaxies to be $b(L_s,z=0) = 1.10 \\pm 0.08$ on comoving scales of $0.02 < k < 0.15 \\hompc$. If luminosity segregation operates on these scales, $L_*$ galaxies would be almost un-biased, $b(L_*,z=0) \\approx 0.96$. These results are derived by assuming a flat $\\Lambda$CDM Universe, and by marginalizing over other free parameters and fixing the spectral index $n=1$ and the optical depth due to reionization $\\tau =0$. We also study the best fit pair $(\\omegam,b)$, and the robustness of the results to varying $n$ and $\\tau$. Various modelling corrections can each change the resulting $b$ by 5--15 per cent. The results are compared with other independent measurements from the 2dFGRS itself, and from the SDSS, cluster abundance and cosmic shear. \\hfill\\break \\hfill\\break {\\bf Key words}: Cosmology, CMB, galaxies, Statistics ", + "introduction": "The 2dF Galaxy Redshift Survey (2dFGRS) has now measured over 210\\,000 galaxy redshifts and is the largest existing galaxy redshift survey (Colless et al. 2001). A sample of this size allows large-scale structure statistics to be measured with very small random errors. Two other 2dFGRS papers, Percival et al. (2001; hereafter P01) and Efstathiou et al. (2001; hereafter E02) have mainly compared the {\\it shape} of the 2dFGRS and CMB power spectra, and concluded that they are consistent with each other (see also Tegmark, Hamilton \\& Xu 2001). Here we estimate the {\\it amplitudes} of the rms fluctuations in mass $\\sigma_{8{\\rm m}}$ and in galaxies $\\sigma_{8{\\rm g}}$. More precisely, we consider the ratio of galaxy to matter power spectra, and use the ratio of these to define the bias parameter: \\begin{equation} b^2 \\equiv {P_{\\rm {gg}}(k)\\over P_{\\rm {mm}}(k)}. \\end{equation} As defined here, $b$ is in principle a function of scale. In practice, we will measure the average value over the range of wavenumbers $0.02 < k < 0.15 \\hompc$. On these scales, the fluctuations are close to the linear regime, and there are good reasons (e.g. Benson et al. 2000) to expect that $b$ should tend to a constant. In this study, we will not test the assumption that the biasing is scale-independent, but we do allow it to be function of luminosity and redshift. A simultaneous analysis of the constraints placed on cosmological parameters by different kinds of data is essential because each probe -- e.g. CMB, Type Ia supernovae (SNe Ia), redshift surveys, cluster abundance, and peculiar velocities -- typically constrains a different combination of parameters (e.g. Bahcall et al. 1999; Bridle et al. 2000, 2001a; E02). A particular case of joint analysis is that of galaxy redshift surveys and the Cosmic Microwave Background (CMB). While the CMB probes the fluctuations in matter, the galaxy redshift surveys measure the perturbations in the light distribution of particular tracer (e.g. galaxies of certain type). Therefore, for a fixed set of cosmological parameters, a combination of the two can tell us about the way galaxies are `biased' relative to the mass fluctuations (e.g. Webster et al. 1998). A well-known problem in estimating cosmological parameters is the degeneracy of parameters, and the choice of free parameters. Here we consider three classes of parameters: (i) Parameters that are fixed by theoretical assumptions or prejudice (which may be supported by observational evidence). Here we assume a flat Universe (i.e. zero curvature), and no tensor component in the CMB (for discussion of the degeneracy with respect to these parameters see E02). (ii) The `free parameters' that are of interest to address a particular question. For the joint 2dFGRS \\& CMB analysis presented here we consider five free parameters: the matter density parameter $\\omegam$, the linear-theory amplitude of the mass fluctuations $\\sigma_{8{\\rm m}}$, the present-epoch linear biasing parameter $b(L_s,z=0)$ (for the survey effective luminosity $L_s \\simeq 1.9 L_*$), the Hubble constant $ h \\equiv H_0/(100 \\kms)$, and the baryon density parameter $\\omega_{\\rm b} \\equiv \\omegab h^2$. As we are mainly interested in combinations of $\\sigma_{8{\\rm m}}$, $b$ and $\\omegam$, we shall marginalize over the remaining parameters. (iii) The robustness of the results to some `extra parameters', that are uncertain. Here we consider the optical depth $\\tau$ due to reionization (see below) and the primordial spectral index $n$. We use as our canonical values $\\tau=0$ and $n=1$, but we also quote the results for other possibly realistic values, $\\tau=(0.05, 0.2)$ and $n=(0.9,1.1)$. The outline of this paper is as follows. In Section 2 we derive $\\sigma_{8{\\rm g}}$ from the 2dFGRS alone, taking into account corrections for redshift-space distortion and for epoch-dependent and luminosity-dependent biasing. In Section 3 we derive $\\sigma_{8{\\rm m}}$ from the latest CMB data. In Section 4 we present a joint analysis of 2dFGRS \\& CMB. Finally, in Section 5 we compare and contrast our measurements with other cosmic probes. ", + "conclusions": "" + }, + "0112/astro-ph0112448_arXiv.txt": { + "abstract": "Using a near-IR stellar library of 706 stars with a wide coverage of atmospheric parameters, we study the behaviour of the Ca\\,{\\sc ii} triplet strength in terms of effective temperature, surface gravity and metallicity. Empirical fitting functions for recently defined line-strength indices, namely CaT$^*$, CaT and PaT, are provided. These functions can be easily implemented into stellar populations models to provide accurate predictions for integrated Ca\\,{\\sc ii} strengths. We also present a thorough study of the various error sources and their relation to the residuals of the derived fitting functions. Finally, the derived functional forms and the behaviour of the predicted Ca\\,{\\sc ii} are compared with those of previous works in the field. ", + "introduction": "This is the third paper in a series dedicated to the understanding of stellar populations of early--type galaxies and other stellar systems by using their near--IR spectra and, particularly, the strength of the integrated Ca\\,{\\sc ii} triplet. In the previous papers, we have presented the basic ingredients of this project. First, we have observed a new stellar library of 706 stars at 1.5~\\AA\\ (FWHM) spectral resolution in the range $\\lambda\\lambda$~8348-9020~\\AA\\ (Cenarro et al. 2001a, Paper~I). That paper also includes the definition of 3 new, improved line-strength indices in the Ca\\,{\\sc ii} triplet region (namely CaT, PaT, CaT$^{*}$), which are especially suited to be measured in the integrated spectra of stellar populations. Paper~II of the series (Cenarro et al. 2001b) presents an updated set of atmospheric parameters ($T_{{\\rm eff}}$, $\\log g$ and [Fe/H]) for the stars of the library. The objective of this third paper is to provide empirical fitting functions describing the behaviour of the above indices in terms of the atmospheric parameters. Together with the spectra of the stellar library, the fitting functions will be implemented into an evolutionary stellar populations synthesis code to predict both the spectral energy distribution and the strengths of the indices for stellar populations of several ages and metallicities (Vazdekis et al. 2001, Paper~IV). A popular method to investigate the star formation history and the stellar content of galaxies is to compare observed line strength indices with stellar population models. This requires an accurate, prior knowledge of the behaviour of the spectral features of interest for a wide range of stellar spectral types, luminosity classes and metallicities, which can be accomplished with the aid of either empirical stellar libraries or theoretical model atmospheres. There are two common alternatives for the stellar populations synthesis approach. The first method mixes the spectra of the different stars with their relative ratios given by evolutionary synthesis models to predict spectral energy distributions (Fioc \\& Rocca-Volmerange 1997; Leitherer et al. 1999; Vazdekis 1999; Schiavon, Barbuy \\& Bruzual 2000; Bruzual \\& Charlot 2001). The major advantage of this spectral synthesis is that it provides full information of all the spectral features within the spectral range covered by the stellar library. However, its usefulness relies on the availability of a complete stellar library at the appropriate spectral resolution. The second procedure, and probably the most widely employed in the past, is the use of empirical fitting functions describing the strength of previously defined spectral features in terms of the main atmospheric parameters. These calibrations are directly implemented into the stellar populations models to derive the index values for populations of different ages and metallicities. It must be noted that one of the main advantages of using fitting functions to reproduce the behaviour of spectral indices is that it allows stellar populations models to include the contribution of all the required stars by means of smooth interpolations between well-populated regions in the parameter space. One should not forget that, in any of the above two methods, a common limitation of the empirical procedures arises from the fact that they implicitly include the chemical enrichment history of the solar neighborhood. This caveat must be kept in mind when using model predictions to interpret the stellar populations of external galaxies, whose star formation histories might be completely different to that of the Galaxy. The usefulness of the fitting functions approach has been clearly demonstrated by the fact that the most important evolutionary synthesis models (e.g., Worthey 1994, Vazdekis et al. 1996, Tantalo et al. 1996 or Bruzual \\& Charlot 2001) have implemented the available fitting functions to fit observed line strengths in the literature. Unfortunately, at present fitting functions are only available in the blue and visible part of the spectrum (e.g. Gorgas et al. 1993, Worthey et al. 1994 and Worthey \\& Ottaviani 1997 for the Lick/IDS indices; Poggianti \\& Barbaro 1997 and Gorgas et al. 1999 for the $\\lambda$ 4000\\AA\\ break). Since other spectral regions also contain very useful, complementary, absorption lines, it is necessary to extend this kind of calibrations to line-strength indices in other spectral regimes, such as the ultraviolet and the near infrared (e.g. the CO index at 2.2~$\\mu$m analyzed by Doyon, Joseph \\& Wright 1994). This is indeed an important motivation for this paper. So, before we give a comprehensive analysis of the predictions of both the spectral synthesis and the fitting functions for stellar populations (Paper IV), we will concentrate in this paper on understanding the behaviour of the Ca\\,{\\sc ii} triplet as a function of the atmospheric parameters, providing the corresponding fitting functions. The Ca\\,{\\sc ii} triplet is one of the most prominent features in the near-IR region of the spectrum of cool stars. Even though there are many previous works dealing with the near-IR Ca\\,{\\sc ii} triplet (we refer the reader to the short review presented in Section~2 of Paper~I), the papers by D\\'{\\i}az, Terlevich \\& Terlevich (1989), Zhou (1991), J$\\o$rgensen, Carlsson \\& Johnson (1992) and Idiart, Thevenin \\& de Freitas-Pacheco (1997) (hereafter DTT, ZHO, JCJ and ITD, respectively) deserve a special mention, since they provide empirical fitting functions for the Ca\\,{\\sc ii} triplet. They have been very useful to obtain a first understanding of the behaviour of the Ca\\,{\\sc ii} triplet as a function of the stellar parameters but, as we will see in this paper, suffer from several limitations that make it impossible for stellar populations models to predict reliable calcium strengths, especially for old stellar populations. Previous papers which have made use of the above mentioned fitting functions to predict the Ca\\,{\\sc ii} triplet in the integrated spectra of galaxies include Vazdekis et al. (1996), ITD, Mayya (1997), Garc\\'{\\i}a--Vargas, Moll\\'{a} \\& Bressan (1998), Leitherer et al. (1999) and Moll\\'{a} \\& Garc\\'{\\i}a--Vargas (2000). Section~2 of this paper describes the qualitative behaviour of the Ca\\,{\\sc ii} triplet as a function of the atmospheric parameters, as derived from the new stellar library. We also include a comparison with the results of the previous work. We devote Section~3 to the mathematical fitting procedure, providing the significant terms, coefficients and statistics of the derived fitting functions. Afterwards, a thorough analysis of residuals and possible error sources is presented, including the sensitivity of the fitting functions to differences in Ca/Fe ratios. In Section~4, we compare the new fitting functions with those presented in previous papers. Finally, Sections~5 is reserved to discuss some important issues and to summarize the contents of this paper. ", + "conclusions": "\\label{conclusions} In this paper, we have analysed the behaviour of the Ca\\,{\\sc ii} triplet strength in the spectra of stars of different spectral types and luminosities, by means of a new near-IR stellar library (presented in Paper~I) with a wide coverage of atmospheric parameters (Paper~II). We have derived empirical fitting functions which can be easily implemented into stellar populations models. Readers interested in including these functions into their population synthesis codes can make use of the {\\sc fortran} subroutine referred to in Section~\\ref{indexdef}. Other results of this work are the following: i) We find a complex behaviour of the Ca\\,{\\sc ii} strength as a function of the three main atmospheric parameters ($T_{\\rm eff}$, $\\log g$ and [Fe/H]). For hot and cold stars, effective temperature and luminosity class are the main driving parameters, whereas, in the mid-temperature regime, all three parameters play an important role (see Fig.~\\ref{fitCaTPaT}). ii) The residuals of the fitting functions arise mainly from uncertainties in the input atmospheric parameters. iii) We do not find any correlation between these residuals and [Ca/Fe] abundance ratios. iv) A comparison with the fitting functions derived in the literature reveals striking differences in the predicted gravity and metallicity dependences for some regions of the parameter space (see Fig.~\\ref{compfit}). Basically, compared to previous fitting functions for the Ca\\,{\\sc ii} triplet, the main advantage of our predictions is that they include the whole range of effective temperatures. This is extremely important since the Ca\\,{\\sc ii} behaviour in the temperature ranges explored by the previous works can not be extrapolated at all to hot ($\\theta \\la 0.75$) or cold stars ($\\theta \\ga 1.35$), as it is readily seen in, for instance, Fig.~\\ref{triplot}. We should not forget that these cold stars constitute a significant contribution to the integrated near-IR spectra of most stellar populations (to old populations through the role of cold dwarfs and giants, and to relatively young populations through the expected contribution of AGB stars; see Paper~IV). This must indeed be remarked since, if the fitting functions by the previous works are extrapolated, stellar populations synthesis models will derive unreliable integrated Ca\\,{\\sc ii} strengths which can not be compared with observations. In the case of synthesis models that make use of the theoretical fitting functions of JCJ, the comparison of the predictions with the observed Ca\\,{\\sc ii} triplet in galactic spectra is even more uncertain since: i) JCJ predictions do not include the contamination of Paschen lines, which affects the measurement of the observed index if the contribution of warm stars is not negligible; and ii) JCJ provides true equivalent widths, whereas line-strength indices in systems like that of DTT are used for the observational data. Although JCJ measure the strengths using the central bandpasses of DTT, these measurements can not be compared to indices in this system, since the local pseudo-continuum (traced between the sidebands) is far from being a true continuum (especially for hot and cold stars; see Paper~I). As a final conclusion, we want to emphasize the importance of using an empirical library with a wide range of accurate atmospheric parameters and a sufficiently large number of calibrating stars for deriving reliable empirical fitting functions. Also, well defined index definitions and objective fitting procedures are a critical factor. To summarize, the functions presented in this paper alleviate these problems and should lead to more reliable comparisons between models and observations than in the past. The first predictions of population synthesis models using these results will be presented in Paper~IV." + }, + "0112/astro-ph0112432_arXiv.txt": { + "abstract": "The energetics of the long duration GRB phenomenon is compared with models of a rotating Black Hole (BH) in a strong magnetic field generated by an accreting torus. A rough estimate of the energy extracted from a rotating BH with the Blandford-Znajek mechanism is obtained with a very simple assumption: an inelastic collision between the rotating BH and the torus. The GRB energy emission is attributed to an high magnetic field that breaks down the vacuum around the BH and gives origin to a e$^\\pm$ fireball. Its subsequent evolution is hypothesised, in analogy with the in-flight decay of an elementary particle, to evolve in two distinct phases. The first one occurs close to the engine and is responsible of energising and collimating the shells. The second one consists of a radiation dominated expansion, which correspondingly accelerates the relativistic photon--particle fluid and ends at the transparency time. This mechanism simply predicts that the observed Lorentz factor is determined by the product of the Lorentz factor of the shell close to the engine and the Lorentz factor derived by the expansion. An anisotropy in the fireball propagation is thus naturally produced, whose degree depends on the bulk Lorentz factor at the end of the collimation phase. ", + "introduction": "At cosmological distances the observed GRB fluxes imply energies of order of up to a solar rest-mass ($\\sim 10^{54}$ erg), and as they vary on timescales of the order of milli seconds from causality arguments these must arise in regions whose size is of the order of kilometres. This implies that an $e^\\pm,\\gamma$ fireball must form, which would expand relativistically. The fireball is energised and possibly collimated, mechanically or magnetically, close to the engine (for reviews see e.g. Piran 1999; Meszaros 2002). Subsequently it adiabatically expands and accelerates, until the Thomson transparency is reached (the opacity being determined by either electron--positron pairs or electrons if the fireball is baryon loaded). The GRB phenomenology -- in particular the fast variability and the detection of $\\gamma$--ray emission from an apparently compact region opaque to electron--positron production via photon--photon interaction -- gives compelling reasons for the bulk motion of the emitting plasma to be highly relativistic with Lorentz factors of the order $\\Gamma\\sim 10^2-10^3$. The degree of isotropy/collimation of the ejected fireball is however still unclear. In fact, as the observer only detects $\\gamma$--ray flux from an angle $\\sim \\Gamma^{-1}$, it is not possible to simply discriminate between an isotropic and a jet-like structure from the observed GRB event. Nevertheless this is in principle possible by adequate sampling and determination of the behaviour of the light curves during the afterglow phase: following the deceleration/sideway expansion of the fireball more and more of the emitting plasma can be seen and a break (and steepening) in the light curve would appear when the whole of the volume becomes observable. Indeed, recently a few GRB afterglows were observed at many wavelengths and suggest an axisymmetric jet-like structure for the fireball, thus strongly reducing the estimate of the energetics with respect to the isotropic case (Frail et al. 2001), although clearly increasing the required GRB rate. The temporal decays of the emission at different frequencies, interpreted according to the fireball model, suggest jet beaming with opening angles $\\theta \\sim 3^{\\circ}$ (Frail et al. 2001). An important inference from these observations is also that the GRB have a typical energy with little intrinsic spread (Frail et al. 2001), although alternative possibilities, such as anisotropy of a collimated fireball, can account for the same observed phenomenology (Zhang \\& Meszaros 2002; Rossi, Lazzati \\& Rees 2002). Found observational trends among timing and spectral properties of GRB as well as numerical results appear also to favour anisotropic distributions of energy/velocity in the fireball (Lloyd-Ronning \\& Ramirez-Ruiz 2002; Salmonson \\& Galama 2002; Zhang, Woosley \\& MacFayden 2002). A further important discovery made by $Beppo$SAX, ASCA and Chandra telescopes, is the presence of iron lines in the X-ray spectra of GRBs (e.g. Amati et al. 2000; Piro et al. 2000; Antonelli et al. 2000). This provides a powerful tool to understand the nature and the environment of GRB primary sources (Vietri et al. 2001; Rees \\& M\\'esz\\'aros 2000): strong iron lines imply a rich environment which may be an argument in favour of massive-star progenitor models of GRB (Woosley 1993, Paczynski 1993; Paczynski 1998; Vietri \\& Stella 1998). These findings have been recently accompanied by the claim of the observation of a complex of soft X--ray lines by XMM-Newton in the spectrum of GRB 011211 (Reeves et al. 2002, see also Watson et al. 2002). They suggest in particular that the high temperature derived from the emitting gas could be interpreted as reheating of pre-ejected material by the GRB itself. These observations are in favour of the interpretation of GRBs as a second step of the residual of the primary explosion (e.g. Vietri \\& Stella 1998): the primary explosion leaves over a compact object that could be a rotating black hole, at the centre of a rarefied atmosphere of ejecta. In such scenario it is plausible that the energy extraction from a rotating BH, through the Blandford-Znajek (BZ) mechanism (Blandford \\& Znajek 1977, Lee et al. 2000), where the external magnetic field can be supplied by a torus circulating around the BH at a distance of the order of the Schwarzchild radius R$_{\\rm s}$. In this paper we focus on two aspects of the 'standard' scenario for the GRB event. The first, developed in Section 2, concerns the extraction of energy from a rotating compact object and its conversion into a photon-e$^{\\pm}$ fireball. Subsequently, in Section 3, we suggest that the acceleration and collimation could occur in two phases, the first one consists in energising and collimating the shells, the second one of a radiation dominated expansion. This mechanism predicts that the observed Lorentz factor is determined by the product of the Lorentz factor of the shell close to the engine and the Lorentz factor derived by the expansion, thus naturally giving rise to an anisotropic fireball. Our conclusions are reported in Section 4. ", + "conclusions": "We considered the possibility that fireballs in long GRB are created by a high magnetic field that breaks down the vacuum around the BH and gives origin to a e$^\\pm$ fireball. The energy can be extracted from a rotating BH via the Blandford-Znajek mechanism thanks to a strong magnetic field generated by an accreting torus. The fireball evolution should then proceed in two phases, the first one consisting in the energisation and collimation of the shells by the external magnetic field and the second one - a radiation dominated expansion - corresponding to the acceleration of the relativistic photon--particle fluid and ending at the transparency radius. This scenario predicts that the resulting Lorentz factor is determined by the product of the Lorentz factor of the shell close to the engine and the Lorentz factor derived by the expansion and simply leads to the formation of an anisotropic fireball. For typical parameters expected in the model the opening angle of the jet obtained in this model could be then estimated to be of order of a few degrees, depending on the efficiency of the acceleration and the resulting angular dependence is similar to what already proposed in the literature on different grounds." + }, + "0112/astro-ph0112551.txt": { + "abstract": "We review the formalism and applications of non-linear perturbation theory (PT) to understanding the large-scale structure of the Universe. We first discuss the dynamics of gravitational instability, from the linear to the non-linear regime. This includes Eulerian and Lagrangian PT, non-linear approximations, and a brief description of numerical simulation techniques. We then cover the basic statistical tools used in cosmology to describe cosmic fields, such as correlations functions in real and Fourier space, probability distribution functions, cumulants and generating functions. In subsequent sections we review the use of PT to make quantitative predictions about these statistics according to initial conditions, including effects of possible non Gaussianity of the primordial fields. Results are illustrated by detailed comparisons of PT predictions with numerical simulations. The last sections deal with applications to observations. First we review in detail practical estimators of statistics in galaxy catalogs and related errors, including traditional approaches and more recent developments. Then, we consider the effects of the bias between the galaxy distribution and the matter distribution, the treatment of redshift distortions in three-dimensional surveys and of projection effects in angular catalogs, and some applications to weak gravitational lensing. We finally review the current observational situation regarding statistics in galaxy catalogs and what the future generation of galaxy surveys promises to deliver. ", + "introduction": "Understanding the large scale structure of the Universe is one of the main goals of cosmology. In the last two decades it has become widely accepted that gravitational instability plays a central role in giving rise to the remarkable structures seen in galaxy surveys. Extracting the wealth of information contained in galaxy clustering to learn about cosmology thus requires a quantitative understanding of the dynamics of gravitational instability and application of sophisticated statistical tools that can best be used to test theoretical models against observations. In this work we review the use of non-linear cosmological perturbation theory (hereafter PT) to accomplish this goal. The usefulness of PT in interpreting results from galaxy surveys is based on the fact that in the gravitational instability scenario density fluctuations become small enough at large scales (the so-called ``weakly non-linear regime'') that a perturbative approach suffices to understand their evolution. Since early developments in the 80's, PT has gone through a period of rapid evolution in the last decade which gave rise to numerous useful results. Given the imminent completion of next-generation large-scale galaxy surveys ideal for applications of PT, it seems timely to provide a comprehensive review of the subject. The purpose of this review is twofold: 1) To summarize the most important theoretical results, which are sometimes rather technical and appeared somewhat scattered in the literature with often fluctuating notation, in a clear, consistent and unified fashion. We tried in particular to unveil approximations that might have been overlooked in the original papers, and to highlight the outstanding theoretical issues that remain to be addressed. 2) To present the state of the art observational knowledge of galaxy clustering with particular emphasis in constraints derived from higher-order statistics on galaxy biasing and primordial non-Gaussianity, and give a rigorous basis for the confrontation of theoretical results with observational data from upcoming galaxy catalogues. We assume throughout this review that the universe satisfies the standard homogeneous and isotropic big bang model. The framework of gravitational instability, in which PT is based, assumes that gravity is the only agent at large scales responsible for the formation of structures in a universe with density fluctuations dominated by dark matter. This assumption is in very good agreement with observations of galaxy clustering, in particular, as we discuss in detail here, from higher-order statistics which are sensitive to the detailed structure of the dynamics responsible for large-scale structures\\footnote{As opposed to just properties of the linearized equations of motion, which can be mimicked by nongravitational theories of structure formation in some cases~\\cite{BWDO94}.}. The non-gravitational effects associated with galaxy formation may alter the distribution of luminous matter compared to that of the underlying dark matter, in particular at small scales: such `galaxy biasing' can be probed with the techniques reviewed in this work. Inevitably, we had to make some decisions in the choice of topics to be covered. Our presentation is definitely focused on the density field, with much less coverage on peculiar velocities. This choice is in particular motivated by the comparatively still preliminary stage of cosmic velocity fields, at least from an observational point of view (see however~\\cite{StWi95,CSW00} for a review). On the other hand, note that since velocity field results are often obtained by identical techniques to those used for the density field, we mention some of these results but without giving them their due importance. In order to fully characterize the density field, we choose to follow the traditional approach of using statistical methods, in particular, $N$-point correlation functions~\\cite{Peebles80}. Alternative methods include morphological descriptors such as Minkowski functionals (of which the genus is perhaps the most widely known), percolation analysis, etc. Unlike correlation functions, however, these other statistics are not as directly linked to dynamics as correlation functions, and thus are not as easy to predict from theoretical models. Furthermore, applications of PT to make predictions of these quantities is still in its infancy (see e.g~\\cite{Matsubara00} and references therein for recent work). Given that PT is an approximate method to solve the dynamics of gravitational clustering, it is desirable to test the validity of the results with other techniques. In particular, we resort to numerical simulations, which involve different approximations in solving the equations of motion that are not restricted to the weakly non-linear regime. There is a strong and healthy interplay between PT and $N$-body simulations which we extensively illustrate throughout this review. At large scales PT can be used to test quantitatively for spurious effects in numerical simulations (e.g. finite volume effects, transients from initial conditions), whereas at smaller, non-linear scales $N$-body simulations can be used to investigate the regime of validity of PT predictions. Although reviewing the current understanding of clustering at small scales is beyond the scope of this review, we have also included a discussion of the predictions of non-linear clustering amplitudes because connections between PT and strongly non-linear behavior have been suggested in the literature. We also include a discussion about stable clustering at small scales which, when coupled with self-similarity, leads to a connection between the large and small-scale scaling behavior of correlations functions. This review is structured so that different chapters can be read independently, although there are inevitable relations. Chapter~2 deals with the basic equations of motion and their solution in PT, including a brief summary of numerical simulations. Chapter~3 is a review of the basics of statistics; we have made it as succinct as possible to swiftly introduce the reader to the core of the review. For a more in-depth treatment we refer the reader to~\\cite{KeSt94,Bertschinger92}. The next two chapters represent the main theoretical results; Chapter~4 deals with $N$-point functions, whereas Chapter~5 reviews results for the smoothed one-point moments and PDF's. These two chapters heavily rely on material covered in Chapters~2 and~3. In Chapter~6 we describe in detail the standard theory of estimators and errors for application to galaxy surveys, with particular attention to the issue of cosmic bias and errors of estimators of the two-point correlation function, power spectrum, and higher-order moments such as the skewness. Chapter~7 deals with theoretical issues related to surveys, such as redshift distortions, projection effects, galaxy biasing and weak gravitational lensing. Chapter~8 presents the current observational status of galaxy clustering, including future prospects in upcoming surveys, with particular emphasis on higher-order statistics. Chapter~9 contains our conclusions and outlook. A number of appendices extend the material in the main text for those interested in carrying out detailed calculations. Finally, to help the reader, Tables~\\ref{tab:abbrev}--\\ref{fb:tabStats} list the main abbreviations and notations used for various cosmological variables, fields and statistics. \\newpage \\begin{table}[p] \\caption{Abbreviations} \\vspace{.1 cm} \\begin{tabular}{cp{12cm }} \\hline PT & Perturbation Theory;\\\\ 2LPT & Second Order Lagrangian Perturbation Theory;\\\\ EPT & Extended Perturbation Theory;\\\\ HEPT & HyperExtended Perturbation Theory;\\\\ ZA & Zel'dovich Approximation;\\\\ SC & Spherical Collapse;\\\\ CDM & Cold Dark Matter (model);\\\\ SCDM & Standard CDM model;\\\\ $\\Lambda$CDM & Flat CDM model with a cosmological constant;\\\\ PDF & Probability Distribution Function;\\\\ CPDF & Count Probability Distribution Function.\\\\ \\hline \\end{tabular} \\label{tab:abbrev} \\end{table} \\begin{table}[p] \\caption{Notation for Various Cosmological Variables} \\vspace{.1 cm} \\begin{tabular}{cp{12cm }} \\hline $\\Omega_m$ & The total matter density in units of critical density;\\\\ $\\Omega_\\Lambda$ & The reduced cosmological constant;\\\\ $\\Omega_{tot}$ & The total energy density of the universe in units of critical density, $\\Omega_{tot}=\\Omega_m+\\Omega_\\Lambda$;\\\\ $H$ & The Hubble constant;\\\\ $h$ & The Hubble constant at present time, in units of $100$ km/s/Mpc, $h \\equiv H_0/100$;\\\\ $a$ & The scale factor;\\\\ $\\tau$ & The conformal time, $d\\tau=dt/a$;\\\\ ${\\cal H}$ & The conformal expansion rate, ${\\cal H}=aH$;\\\\ $D_1$ & The linear growth factor;\\\\ $D_n$ & The $n$-th order growth factor;\\\\ $f(\\Omega_m,\\Omega_\\Lambda)$ & The logarithmic derivative of (the fastest growing mode of) the linear growth factor with respect to $a$: $f(\\Omega_m,\\Omega_\\Lambda)\\equiv d \\ln D_1/d\\ln a$.\\\\ \\hline \\end{tabular} \\label{tab:cosmopar} \\end{table} \\begin{table}[p] \\caption{Notation for the Cosmic Fields} \\vspace{.1 cm} \\begin{tabular}{cp{12cm }} \\hline ${\\tilde X}$ & The Fourier transform of field $X$;\\\\ & $\\tilde{X}(\\vk)=(2\\pi)^{-3} \\int \\d^3\\vx \\,{\\rm e}^{-\\ii\\vk\\cdot\\vx}\\,X(\\vx)$ (except in Sect.~\\ref{sec:sec3}) \\\\ $\\vx$ & The comoving position in real space;\\\\ $\\rho(\\vx)$ & The local cosmic density;\\\\ $\\delta(\\vx)$ & The local density contrast, $\\delta=\\rho/\\rhobar-1$;\\\\ $\\Phi(\\vx)$ & The gravitational potential;\\\\ $\\vu(\\vx)$ & The local peculiar velocity field;\\\\ $\\theta(\\vx)$ & The local velocity divergence in units of ${\\cal H}=aH$;\\\\ $F_p(\\vk_1,\\dots,\\vk_p)$ & The $p^{\\rm th}$ order density field kernel;\\\\ $G_p(\\vk_1,\\dots,\\vk_p)$ & The $p^{\\rm th}$ order velocity divergence field kernel;\\\\ $\\psi(\\vq)$ & The Lagrangian displacement field;\\\\ $J(\\vq)$ & The Jacobian of the Lagrangian-Eulerian mapping.\\\\ \\hline \\end{tabular} \\label{fb:tabFields} \\end{table} { } \\begin{table}[p] \\caption{Notation for Statistical Quantities} \\vspace{.1 cm} \\begin{tabular}{cp{12cm }} \\hline $P(k)$ & The density power spectrum;\\\\ $\\Delta(k)$ & The dimensionless power, $\\Delta=4\\pi k^3 P(k)$;\\\\ $B(k_1,k_2,k_3)$ & The bispectrum;\\\\ $P_N(\\vk_1,\\ldots,\\vk_N)$ & The $N$-point polyspectrum;\\\\ $P_N$ & The count-in-cell probability distribution function;\\\\ $p(\\delta)\\d\\delta$ & The cosmic density probability distribution function;\\\\ $F_k$ & The factorial moment of order $k$;\\\\ $\\xi_2(\\vx_1,\\vx_2)\\equiv \\xi_{12}\\equiv \\xi$ & The two-point correlation function, $\\xi_2(\\vx_1,\\vx_2)=\\langle \\delta(\\vx_1) \\delta(\\vx_2) \\rangle=\\langle \\delta(\\vx_1) \\delta(\\vx_2) \\rangle_c$;\\\\ $\\sigma^2\\equiv \\overline{\\xi}\\equiv \\overline{\\xi}_2$ & The cell-average $two-$point correlation function;\\\\ $\\sigma_8$ & The value of the (linearly extrapolated) $\\sigma$ in a sphere of $8\\,h^{-1}$ Mpc radius;\\\\ $\\Gamma$ & Shape parameter of the linear power-spectrum, $\\Gamma \\simeq \\Omega_m h$;\\\\ \\hline \\end{tabular} \\label{fb:tabStats} \\end{table} \\begin{table} \\noindent Table~\\ref{fb:tabStats} (continued) \\vspace{.1 cm} \\begin{tabular}{cp{12cm }} \\hline $\\xi_N(\\vx_1,\\dots,\\vx_N)$ & The $N$-point correlation functions $\\xi_N(\\vx_1,\\dots,\\vx_N)=\\langle \\delta(\\vx_1)\\ldots\\delta(\\vx_N) \\rangle_c$;\\\\ $w_N(\\theta_1,\\ldots,\\theta_N)$ & The angular $N$-point correlation functions;\\\\ $\\overline{\\xi}_N$ & The cell-averaged $N-$point correlation functions $\\overline{\\xi}_N=\\langle \\delta_R^N \\rangle_c$;\\\\ $\\overline{w}_N$ & The cell-averaged angular $N-$point correlation functions;\\\\ $S_p$ & The density normalized cumulants, $S_p=\\langle\\delta_R^p\\rangle_c/\\langle \\delta_R^2 \\rangle^{p-1} =\\overline{\\xi}_p/\\overline{\\xi}^{p-1}$;\\\\ $S_3$, $S_4$ & The (reduced) skewness/kurtosis;\\\\ $s_p$ & The projected density normalized cumulants;\\\\ $Q\\equiv Q_3$, ${\\tilde Q}\\equiv{\\tilde Q}_3$ & The three-point hierarchical amplitude in real/Fourier space;\\\\ $Q_N$, $\\tilde{Q}_N$ & The $N$-point hierarchical amplitude in real/Fourier space; $Q_N$ can also stand for $S_N/N^{N-2}$ (Chap.~\\ref{sec:chapter7});\\\\ $q_N$, $\\tilde{q}_N$ & The projected $N$-point hierarchical amplitude in real/Fourier space; $q_N$ can also stand for $s_N/N^{N-2}$ (Chap.~\\ref{sec:chapter7});\\\\ $T_p$ & The velocity divergence normalized cumulants;\\\\ $C_{pq}$ & The~two-point~density~normalized~cumulants, $C_{pq}= \\langle\\delta_1^p \\delta_2^q \\rangle_c/(\\xi_{12}\\mom2^{p+q-2})$;\\\\ $\\varphi(y)$ & The one-point cumulant generating function, $\\varphi(y)=\\sum_pS_p\\,(-y)^p/p!$;\\\\ $\\nu_p$, $\\mu_p$ & The density/velocity field vertices;\\\\ $\\mGd(\\tau)\\equiv\\mGd^L(\\tau) $, $\\mGv(\\tau)\\equiv\\mGv^L(\\tau) $ & The vertex generating function for the density/velocity field, $\\mGd(\\tau)\\equiv \\sum_{p\\geq 1} \\nu_p (-\\tau)^p/p!$, and $\\mGv(\\tau)\\equiv -f(\\Omega_m,\\Omega_\\Lambda)\\sum_{p\\geq 1} \\nu_p (-\\tau)^p/p!$;\\\\ $\\langle X \\rangle$ & The ensemble average of statistic $X$;\\\\ ${\\hat X}$ & The estimator of statistic $X$;\\\\ $\\Upsilon(\\hat X)\\d{\\hat X}$ & The cosmic distribution function of estimator ${\\hat X}$;\\\\ $\\Delta X$ & The cosmic error on estimator ${\\hat X}$.\\\\ \\hline \\end{tabular} \\end{table} \\clearpage %%%%%%%%%%%%%%%%%%%%%%%%%%%% CHAPTER2 %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ", + "conclusions": "As illustrated throughout this work, PT provides a valuable tool to understand and calculate predictions for the evolution of large scale structure in the universe. The last decade has witnessed a substantial activity in this area, with strong interplay with numerical simulations of structure formation and observations of clustering of galaxies and, more recently, weak gravitational lensing. As galaxy surveys become larger probing more volume in the weakly non-linear regime, new applications of PT are likely to flourish to provide new ways of learning about cosmology, the origin of primordial fluctuations, and the relation between galaxies and dark matter. The general framework of these calculations is well established and calculations have been pursued for a number of observational situations, whether it is for the statistical properties of the local density contrast, the velocity divergence, for the projected density contrast, redshift measurements or for more elaborate statistics such as joint density cumulants. All these results provide robust frameworks for understanding the observations or for reliable error computations. There are however a number of outstanding issues that remain to be addressed in order to improve our understanding of gravitational instability at large scales, \\\\ \\begin{itemize} \\item[--] Most of the calculations have been done assuming Gaussian initial conditions, except for some specific cases such as $\\chi^{2}$ models. Although present observations are consistent with Gaussian initial conditions, deriving quantitative constraints on primordial non-Gaussianity requires some knowledge or useful parametrization of non-Gaussian initial conditions and how they evolve by gravity. \\item[--] Predictions of PT for velocity field statistics are still in a rudimentary state compared to the case of the density field. Upcoming velocity surveys will start probing scales where PT predictions can be used. In addition, robust methods for calculating redshift distortions including the non-linear effects due to the redshift-space mapping are needed to fully extract information from the next generation of galaxy redshift surveys. \\item[--] Another observational context in which a PT approach can be very valuable is the Lyman-$\\alpha$ forest observed in quasar spectra. The statistical properties of these systems should be accessible to perturbative methods since most of the absorption lines correspond to modest density contrasts (from 1 to 10). This is a very promising field for observational cosmology. \\item[--] Accurate constraints on cosmological parameters from galaxy surveys require precise models of the joint likelihood of low and higher-order statistics including their covariance matrices. To date this has only been investigated in detail numerically, or analytically in some restricted cases. \\\\ \\end{itemize} In addition, as we probe the transition to the non-linear regime, there are a few technical issues that need more investigation, \\\\ \\begin{itemize} \\item[--] Most results have been obtained in the tree-level approximation, for which systematic calculations can be done and the emergence of non-Gaussianity can be characterized in an elegant way. There is no such systematic framework for loop corrections, and only a few general results are known in this case. Furthermore, loop corrections are found to be divergent for power-law spectra with index $n \\geq -1$, the interpretation of which is still not clear. Although this issue is irrelevant for realistic spectra such as CDM, its resolution may shed some light into the physics of the non-linear regime. \\item[--] The SC collapse prescription (Sect.~\\ref{sec:SCmodel}) leads to a good description of $S_{p}$ parameters in the transition to the non-linear regime when compared to N-body simulations and exact one-loop corrections when known. Is it possible to improve on this approximation, or make it more rigorous in any well-controlled way while maintaining its simplicity? \\item[--] The development of HEPT (Sect.~\\ref{sec:HEPT}) and EPT (Sect.~\\ref{sec:EPT}) suggests that there is a deep connection between gravitational clustering at large and small scales. Is this really so, or is it just an accident? Why do strongly non-linear clustering amplitudes seem to be so directly related to initial conditions? \\\\ \\end{itemize} {}From the observational point of view, the next few years promise to be extremely exciting, with the completion of 2dFGRS and SDSS and deep surveys that will trace the evolution of large-scale structure towards high redshift\\footnote{~See e.g.~\\cite{CDS01} for a recent assessment of how well upcoming deep surveys will determine correlation functions.}. Observations of the so-called Lyman break galaxies~\\cite{SADGPK98} are should soon provide a precious probe of the high-redshift universe, in particular regarding the evolution of galaxy bias~\\cite{ASGDPK98,PoGi01,BWS02}. Furthermore, weak lensing observations will provide measurements of the projected mass density that can be directly compared with theoretical predictions. In addition, CMB satellites and high-resolution experiments will probe scales that overlap with galaxy surveys and thus provide a consistency check on the framework of the growth of structure. Outstanding observational issues abound, most of them perhaps related to the way galaxies form and evolve. One of the most pressing ones, as discussed many times in Chapter~8, is probably to have a convincing explanation of why correlation functions scale as power-law's at non-linear scales. The scaling in Figs.~\\ref{w2apmz4} and~\\ref{PkPSCz} is certainly remarkable and preliminary results from 2dFGRS~\\cite{NBHMPCF01} and SDSS~\\cite{Zehavi01SDSS,Connolly01SDSS} seem already to confirm and extend these results. In the CDM framework, however, this simple behavior is thought to be the result of accidental cancellation of the dark matter non-power-law form by scale-dependent bias due to the way dark matter halos are populated by galaxies (see discussion in Sect.~\\ref{dhbias}). Although this may seem rather adhoc, this model has, on the other hand, many observable consequences. The same weighting that makes the two-point function depend as a power-law of separation~\\cite{Seljak00,PeSm00,SSHJ01} suppresses the velocity dispersion and mean streaming of galaxies~\\cite{SHDS01,SDHS01} as observed, see e.g.~\\cite{JMB98}. In addition, this weighting affects higher-order statistics in the non-linear regime, suppressing them in comparison with their dark matter counterparts~\\cite{SSHJ01} (see Fig.~\\ref{Sp_gal}) as observed, see e.g. Fig.~\\ref{s34apm} for a comparison between dark matter and $S_{p}$ in the APM survey. There are also complementary indications that galaxies do not trace the underlying dark matter distribution at small scales from measurements of higher-order statistics. As discussed in Sect.~\\ref{sec:w2pk}, reconstruction of the linear power spectrum from galaxy surveys leads to significant disagreement of higher-order moments if no biasing is imposed at small scales, as shown in Fig.~\\ref{s34apm} for APM galaxies. A promising way to confirm that the underlying higher-order statistics of the dark matter are much higher than those of galaxies at small scales is by measuring higher-order moments in weak gravitational lensing. This will likely be done in the near future, as weak lensing surveys are already beginning to probe the relation between dark matter and galaxies~\\cite{HYG01}. In any case, statistical analysis of future observations are going to decide whether the small-scale behavior of correlations is dictated by biasing or if a new framework is needed to understand galaxy clustering at non-linear scales. What seems clear, whatever the outcome, is that the techniques described here will be a valuable tool to achieve that goal. \\vskip 2pc This project was possible thanks to the hospitality of several institutions that supported frequent visits. We thank CSIC, IAP, IAS, IEEC, SPhT, and also CITA during the initial stages of this work. We also benefited greatly from discussions with F. Bouchet, J. Frieman, J. Fry, R. Juszkiewicz and I. Szapudi that help set the structure of this review. We thank Marc Kamionkowski for many helpful comments about a previous version of this work. This project has made extensive use of NASA's Astrophysics Data System Abstract Service. \\clearpage %%%%%%%%%%%%%%%%%%%%%%%%%%%% APPENDICES %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\appendix" + }, + "0112/astro-ph0112268_arXiv.txt": { + "abstract": "We describe the star formation histories of three late-type dwarfs located outside the Local Group: IZw18, NGC1569 and NGC1705. The results are based on the application of the method of synthetic colour-magnitude diagrams to deep HST photometric data. All the examined galaxies were already forming stars at the reached lookback time and show no evidence of long quiescent phases. The obtained scenarios are quite similar to those derived for other galaxies of this morphological type, both inside and outside the Local Group. ", + "introduction": "The application of the synthetic colour-magnitude diagram (CMD) method to dwarf galaxies in the Local Group has allowed several people to infer the star formation histories (SFHs) of these galaxies with unprecedented accuracy (e.g. Tosi et al. 1991, Aparicio et al. 1996, Tolstoy \\& Saha 1996, Grebel 1998). It has thus been possible to understand that late-type dwarfs evolve following a {\\it gasping} (Tosi et al 1991) regime of star formation (SF) rather than a {\\it bursting} one. In other words, their SF occurs in long episodes of moderate activity, separated by short quiescent phases, and not in short episodes of strong activity, separated by long quiescent periods. The natural question is what do late-type galaxies outside the Local Group do, and what is the SF regime in Blue Compact Dwarfs (BCDs), which are not present locally. BCDs have been suggested to undergo a few strong and short bursts of SF, to be so poorly evolved to represent the closest analogues to primeval galaxies and to be possible contributors to the excess of faint blue objects found in deep galaxy counts at redshifts between 0.7 and 1. Hence, understanding whether or not their SF activity can be strong enough to allow for sufficient brightness at intermediate redshift and whether or not any of them are currently forming their first stars can have interesting cosmological implications. In order to answer these questions, we are studying the SFHs of a number of late-type dwarfs (both BCDs and dwarf irregulars) outside the Local Group. Their individual stars have been resolved by deep Hubble Space Telescope (HST) photometry and their stellar populations can be interpreted in terms of SFH with the synthetic CMD method. So far, we have examined three prototype systems: IZw18, NGC1569 and NGC1705. IZw18 (at a distance between 10 and 14 Mpc) is the most metal poor galaxy ever observed and has often been suggested to be experiencing now its first burst of SF. NGC1569 (at 2.2 Mpc) is one of the most active starburst dwarfs, exhibiting three super star clusters, a large concentration of giant HII regions, and shells and filaments presumably related to young SN ejecta. NGC1705 (at 5.1 Mpc) is a post-starburst BCD, containing a super star cluster and showing the best observational evidence of gas outflows (galactic winds) triggered by SN explosions. ", + "conclusions": "" + }, + "0112/astro-ph0112097_arXiv.txt": { + "abstract": "We report X-ray imaging - spectroscopy of the jet of M87 at sub arc second resolution with the Chandra X-ray Observatory. The galaxy nucleus and all the knots seen at radio and optical wavelengths, as far from the nucleus as knot C, are detected in the X-ray observations. There is a strong trend for the ratio of X-ray to radio, or optical, flux to decline with increasing distance from the nucleus. At least three knots are displaced from their radio/optical counterparts, being tens of pc closer to the nucleus at X-ray than at radio or optical wavelengths. The X-ray spectra of the nucleus and knots are well described by power laws absorbed by cold gas, with only the unresolved nucleus exhibiting intrinsic absorption. In view of the similar spectra of the nucleus and jet knots, and the high X-ray flux of the knots closest to the nucleus, we suggest that the X-ray emission coincident with the nucleus may actually originate from the pc -- or sub-pc -- scale jet rather than the accretion disk. Arguments are given that the X-ray emission process is unlikely to be inverse Compton scattering. Instead, we favor synchrotron radiation. Plotted as $\\nu$S$_{\\rm \\nu}$, the spectra of the knots generally peak in or just above the optical - near infrared band. However, the overall spectra of at least three knots cannot be described by simple models in which the spectral index monotonically increases with frequency, as would result from synchrotron losses or a high energy cut-off to the injected electron spectrum. Instead, these spectra must turn down just above the optical band and then flatten in the X-ray band. In the context of a synchrotron model, this result suggests that either the X-ray emitting electrons/positrons in these knots represent a separate ``population'' from those that emit the radio and optical radiation or the magnetic field is highly inhomogeneous. If the former interpretation is correct, our results provide further support for the notion that radio galaxies produce a hard ($\\gamma$ $\\simeq$ 2 -- 2.5, N(E) $\\propto$ E$^{-\\gamma}$) spectrum of high energy (E/m$_{e}$c$^{2}$) $\\sim$ 10$^{7-8}$) electrons and possibly positrons. ", + "introduction": "The jet in M87 is arguably the most famous of its kind, primarily because of its proximity and strong optical synchrotron radiation. It has been studied in great detail at radio (see Biretta 1999 for a review) and optical (e.g. Biretta, Sparks \\& Macchetto 1999; Perlman et al. 2001) wavelengths. Radio observations of the jet show that it is polar down to 0.01 pc (Junor \\& Biretta 1995), which is only 60 times the gravitational radius of the black hole (mass $\\simeq$ 3 $\\times$ 10$^{9}$ M$_{\\odot}$, e.g. Macchetto et al. 1997). HST observations have revealed numerous features in the inner $\\simeq$ 400 pc with speeds in the range 4c to 6c, confirming that the bulk flow of the jet is relativistic (Biretta, Sparks \\& Macchetto 1999). At X-ray wavelengths, the nucleus and knot A have been detected (Schreier, Gorenstein \\& Feigelson 1982; Biretta, Stern \\& Harris 1991; Neumann et al. 1997; Harris, Biretta \\& Junor 1997; B\\\"ohringer et al. 2001). The limited resolution of prior X-ray telescopes has precluded a study with comparable resolution to optical telescopes. With the advent of the Chandra X-ray Observatory, it is now possible to obtain imaging - spectroscopy of the jet with sub arc second resolution. In Section 2, we describe the Chandra observations and the data reduction. Section 3 presents the results, discussing the morphology of the jet and the spectra of the nucleus and individual jet knots. Section 4 is devoted to the radiation mechanism of the X-rays, while Section 5 summarises our conclusions. We assume a distance of 16 Mpc (Tonry 1991; Whitmore et al. 1995), so 1$^{\\prime\\prime}$ = 78 pc. The Galactic column density in the direction of M87 is N$_{\\rm H}$ (Gal) = 2.5 $\\times$ 10$^{20}$ cm$^{-2}$ (Stark et al. 1992). ", + "conclusions": "We have obtained high sensitivity, X-ray imaging - spectroscopy of the jet of M87 with sub arc second angular resolution. Superficially, the X-ray jet resembles that in the radio and optical bands, with all knots out to knot C detected. However, there is a very strong trend for the ratio of X-ray to either radio or optical flux to decline with increasing distance from the nucleus. This strength of the near-nuclear jet in X-rays suggests that the X-ray emission coincident with the nucleus may actually originate from the pc -- or sub-pc -- scale jet, rather than the accretion disk. There are clear morphological differences between the radio/optical jet, on the one hand, and the X-ray jet, on the other. In particular, some knots in the X-ray image are displaced from their radio/optical counterparts towards the nucleus by tens of pc. The spectra of the nucleus and all jet knots may be well described by power-law spectra absorbed by cold matter. Only the nucleus shows clear evidence for intrinsic absorption, with an equivalent hydrogen column density of $\\sim$ (3 -- 5) $\\times$ 10$^{20}$ cm$^{-2}$. The knots in the jet have photon indices, $\\Gamma$, in the range 2.04 to 2.90, with average $\\Gamma$ $\\simeq$ 2.4. The X-ray spectra of the jet knots are thus considerably steeper than those at radio or optical wavelengths. Plotted as $\\nu$S$_{\\nu}$, the spectra of the knots peak in or somewhat above the optical - near infrared band. The X-ray photon index of the nucleus is $\\Gamma$ = 2.2$\\pm$0.1, similar to the jet knots and again suggesting that the nuclear X-ray emission originates from the jet. We have discussed the process responsible for the X-ray emission. Synchrotron self-Compton emission falls short by a factor of $\\sim$ 100 -- 1,000 for an equipartition field. The field needs to be $\\sim$ 70 times below equipartition to match the X-ray flux. Models invoking inverse Compton scattering of the microwave background require values of $\\delta$ in the range 10 - 40 and an implausibly small angle ($\\theta$ = 1$^{\\circ}$ - 5$^{\\circ}$) between the jet and the line of sight (Harris \\& Krawczynski 2002). Furthermore, the spectral indices of the X-ray and radio emissions are very different, in contrast with the similar spectral indices expected if the synchrotron-emitting and inverse Compton scattering electron populations are the same. For these reasons, we consider all inverse Compton models to be extremely implausible. The X-ray emission from the jet is almost certainly synchrotron radiation. However, for at least three knots, the X-ray spectrum is not a simple continuation of the radio -- optical spectrum. Instead, the spectrum must turn down at frequencies above the optical -- near infrared band, and then flatten in X-rays. The broad-band spectra of these knots are remarkably similar to that of the western hot spot of Pictor A. The optical, near infrared and X-ray data were not taken simultaneously, raising the possibility that our broad-band spectra may be affected by variability. However, the magnitude of the knot variability required to reconcile the observed spectra with simple synchrotron models, in which the spectral index increases monotonically with increasing frequency, is much larger than any yet observed. We briefly discuss the potential application to M87 of a model, originally due to Burn (1973), in which the magnetic field is extremely inhomogeneous. In this model, the X-ray emission originates from low energy electrons temporarily in regions of very strong magnetic field; the X-ray spectrum is related to the spectrum of the strength of the magnetic field rather than the energy spectrum of the relativistic electons. It is, however, very doubtful whether the large field filaments (with magnetic fields enhanced by a factor of $\\sim$ 100) required can be created in the jet plasma. A third alternative is that the X-ray synchrotron-emitting electrons/positrons are a separate ``population'' to those that emit the radio and optical synchrotron radiation (i.e. the energy spectrum and normalisation of the X-ray emitting electrons at injection is not a simple continuation of that of the radio- and optical-emitting electrons). If the last alternative is correct, our X-ray spectra of the M87 jet provide additional support for the notion that radio galaxies produce a hard ($\\gamma$ = 2 - 2.5) spectrum of high energy (E/m$_{\\rm e}$c$^{2}$ $\\sim$ 10$^{7-8}$) electrons and possibly positrons. We are grateful to John Biretta and Ismael P\\'erez-Fournon for providing a 6 cm VLA map and a V band image of the jet, respectively, and to Carole Mundell for making a 2 cm map from the VLA archive. We thank the staff of the Chandra Science Center, especially D. E. Harris and S. Virani, for their help, C. D. Dermer for a valuable correspondence, E. S. Perlman for helpful comments and D. E. Harris for pointing out an error in an earlier version. This research was supported by NASA through grants NAG 81027 and NAG 81755." + }, + "0112/hep-ph0112113_arXiv.txt": { + "abstract": " ", + "introduction": "A favoured candidate for cold dark matter is the lightest supersymmetric particle (LSP), which is generally thought to be the lightest neutralino $\\ch$~\\cite{EHNOS} in the minimal supersymmetric extension of the Standard Model (MSSM). It is common to focus attention on the constrained MSSM (CMSSM), in which all the soft supersymmetry-breaking scalar masses $m_0$ are required to be equal at an input superysmmetric GUT scale, as are the gaugino masses $m_{1/2}$ and the trilinear soft supersymmetry-breaking parameters $A_0$. These assumptions yield well-defined relations between the various sparticle masses, and correspondingly more definite predictions for the relic abundance $\\Omega_\\ch h^2$ and observable signatures. This paper is devoted to relic-abundance calculations including coannihilations of the lightest neutralino $\\ch$ with ${\\tilde t}_1$, the lighter supersymmetric partner of the top quark~\\cite{stopco}. The range $0.1 < \\Omega_\\ch h^2 < 0.3$ is generally thought to be preferred by astrophysics and cosmology~\\cite{omegah2}. Lower values of $\\Omega_\\ch h^2$ might be possible if there is some other source of cold dark matter, but higher values are incompatible with observation. The regions of the $m_{1/2}, m_0$ plane where the relic density falls within the preferred range $0.1 < \\Omega_\\ch h^2 < 0.3$ have generally been divided into four generic parts. There is a `bulk' region at moderate $m_{1/2}$ and $m_0$~\\cite{EHNOS}. Then, extending to larger $m_{1/2}$, there is a `tail' of the parameter space where the LSP $\\ch$ is almost degenerate with the next-to-lightest supersymmetric particle (NLSP), which is in this region the ${\\tilde \\tau}_1$, the lighter supersymmetric partner of the $\\tau$ lepton. Along this `tail', efficient coannihilations~\\cite{gs,oldcoann,eg} keep $\\Omega_\\ch h^2$ down in the preferred range, even for larger values of $m_\\ch$~\\cite{efo,efosi,coann,ADS}. At larger $m_0$, close to the boundary where electroweak symmetry breaking is no longer possible, there is the `focus-point' region where the LSP has a larger Higgsino component and $m_\\ch$ is small enough for $\\Omega_\\ch h^2$ to be acceptable~\\cite{focus}. Finally, extending to larger $m_{1/2}$ and $m_0$ at intermediate values of $m_{1/2} / m_0$, there may be a `funnel' of CMSSM parameter space where rapid direct-channel annihilations via the poles of the heavier Higgs bosons $A$ and $H$ keep $\\Omega_\\ch h^2$ in the preferred range~\\cite{EFGOSi,funnel}. In this paper, we emphasize the significance of coannihilation of the LSP $\\ch$ with ${\\tilde t}_1$, the lighter supersymmetric partner of the $t$ quark \\cite{stopco}. This mechanism opens up another `tail' of parameter space, this time extending to larger values of $m_0$. It is not relevant for the small values of $A_0$ considered in previous coannihilation calculations~\\cite{efosi,coann,ADS}, but may be important for large $A_0$, as we demonstrate in this paper. Coannihilations of $\\ch$ with ${\\tilde t}_1$ are important when the latter is the NLSP, just as $\\ch {\\tilde \\tau}_1$ coannihilations are important when the ${\\tilde \\tau}_1$ is the NLSP. In the latter case, one must also consider coannihilations with the ${\\tilde e}_1$ and ${\\tilde \\mu}_1$, which are not much heavier than the ${\\tilde \\tau}_1$~\\cite{efo,efosi,coann,ADS}. There are also regions of CMSSM parameter space where both the ${\\tilde t}_1$ and ${\\tilde \\tau}_1$ are close in mass to the LSP $\\ch$, and ${\\tilde t}_1 {\\tilde \\tau}_1$ coannihilations must also be considered. We present here detailed calculations of the matrix elements and cross sections for all the leading $\\ch {\\tilde t}_1$ and ${\\tilde t}_1 {\\tilde \\ell}$ coannihilation processes, and illustrate their importance for $\\Omega_\\ch h^2$ in some instances in the CMSSM when $A_0 \\ne 0$. The structure of the paper is as follows. In Section 2 we recall some important features of LSP relic-density calculations in general, and coannihilations in particular. Then, in Section 3 we compare the relative magnitudes of the $\\ch \\ch$, $\\ch {\\tilde t}_1$, ${\\tilde t}_1 {\\tilde t}^{(*)}_1$ and ${\\tilde t}_1 {\\tilde \\ell}^{(*)}$ processes for some specific choices of the CMSSM parameters. Section 4 provides an overview of the implications of $\\ch {\\tilde t}_1$ coannihilation and related processes for the regions of the $m_{1/2}, m_0$ plane allowed by the constraint $0.1 < \\Omega_\\ch h^2 < 0.3$ for various choices of the other CMSSM parameters. Relevant details of our calculations of the matrix elements are contained in an Appendix. ", + "conclusions": "We have documented in this paper the potential importance of $\\ch {\\tilde t_1}$ coannihilation in delineating the preferred domain of CMSSM parameter space for $A_0 \\ne 0$. In this paper, we have only sctratched the surface of this subject, whose higher-dimensional parameter space merits more detailed exploration. The Appendix provides details of the diagrammatic calculations that should be sufficient for our results to be verified and used by other authors. Although applied in the context of the CMSSM, our results may also be used in more general MSSM contexts. However, other coannihilation processes are also important in other regions of the general MSSM parameter space. For example, in the CMSSM the sbottom mass is generally larger than the stop mass even for large $\\tan \\beta$. However, if one allow non-degeneracy in the scalar soft breaking mass term, a sbottom NLSP becomes possible~\\cite{ADS}. A complete calculation of the LSP relic density in the MSSM requires a careful discussion of all such coannihilation possibilities. \\ \\vskip 0.5in \\vbox{ \\noindent{ {\\bf Acknowledgments} } \\\\ \\noindent We thank Toby Falk for many related discussions. The work of K.A.O. and Y.S. was supported in part by DOE grant DE--FG02--94ER--40823.} \\newpage \\input APPENDIX" + }, + "0112/astro-ph0112083_arXiv.txt": { + "abstract": "Gamma ray bursts are often modelled as jet-like outflows directed towards the observer; the cone angle of the jet is then commonly inferred from the time at which there is a steepening in the power-law decay of the afterglow. We consider an alternative model in which the jet has a beam pattern where the luminosity per unit solid angle (and perhaps also the initial Lorentz factor) decreases smoothly away from the axis, rather than having a well-defined cone angle within which the flow is uniform. We show that the break in the afterglow light curve then occurs at a time that depends on the viewing angle. Instead of implying a range of intrinsically different jets -- some very narrow, and others with similar power spread over a wider cone -- the data on afterglow breaks could be consistent with a standardized jet, viewed from different angles. We discuss the implication of this model for the luminosity function. ", + "introduction": "There are strong reasons for suspecting that the emitting plasma of $\\gamma$-ray bursts (GRBs) is geometrically beamed in a cone. The energy requirements can then be reduced below the exorbitant levels that isotropic emission would imply (Kulkarni et al. 1999) and in most models for the long bursts it is in any case a natural expectation -- borne out by simulations (MacFadyen \\& Woosley 1999; MacFadyen, Woosley \\& Heger 2001) -- that the relativistic outflow from a central engine should be collimated along a channel that opens up along the rotation axis of a massive star. Although the relativistic MHD that gives rise to jets is uncertain, and surely very complicated, most discussions of the radiation from gamma ray bursts (and their afterglows) has postulated a jet with a well-defined angle (Meszaros \\& Rees 1997), though this angle may differ for different bursts (see, however, Meszaros, Rees \\& Wijers 1998; Salmonson 2001). We discuss here an alternative model where the jet, rather than having a uniform profile out to some definite cone angle, has a ``beam pattern'' where the power per unit solid angle (and perhaps also the initial Lorentz factor) is maximal along the axis, but drops off gradually away from the axis. This would be expected if there is mixing and entrainment from the borders of the funnel. We discuss the expected time-dependence of the afterglow if it is triggered by a jet with this more general profile. We conclude that the time of the observed break (which, for a uniform jet viewed along its axis, depends on the cone angle; see, e.g., Rhoads 1997) instead depends on the angle between the line of sight and the symmetry axis: such a jet viewed nearly head-on simulates a narrow uniform jet, whereas the afterglow from the same jet viewed more obliquely would simulate a wider uniform jet. GRB have been proposed recently to be explosions releasing a standard power that can be either injected in very different jet opening angles or distributed within the jet in some universal emission diagram, (Postnov et al. 2001). While Frail et al. (2001, hereafter F01), support with their data the first interpretation, we show that their observational results could instead be attributed to a more standard set of objects viewed at different angles to their symmetry axis. ", + "conclusions": "We considered inhomogeneous GRBs jets with a standard total energy, opening angle and local energy distribution, $\\epsilon\\propto\\theta ^{-2}$. We show that this jet structure can reproduce the observed correlation between isotropic energy and break-time. In this model both measurements depend only on the viewing angle because the $\\gamma$-ray fluence and the afterglow emission are dominated by the components of the jet pointing towards the observer at small times. Since all cones have the same total energy $E_{\\rm jet}=2\\pi\\theta^{2}\\epsilon\\,=cost$, we recover the results of F01 and PK01 but the constrains on the geometrical beaming can be relaxed and an appealing more standard structure for all GRBs can be adopted. The jet total energy can be calculated from Eq.~\\ref{eq:top} \\begin{equation} E_{Total}=2\\pi\\epsilon_{c} \\theta_{c}^{2}\\,\\,\\left(1+2 \\ln\\frac{\\theta_{j}}{\\theta_{c}}\\right) = E_{\\rm jet} \\,\\left(1+2 \\ln\\frac{\\theta_{j}}{\\theta_{c}}\\right) \\end{equation} and compared to $E_{\\rm jet}=2\\pi\\theta_{o}\\epsilon_{o}=2\\pi\\epsilon_{c} \\theta_{c}^{2}$, the total energy inferred from observation, $E_{\\rm jet} \\leq E_{Total}$. To give an example, for a fireball with $\\theta_c=1^\\circ$ and $\\theta_j=20^\\circ$, we have $E_{Total}/E_{\\rm jet}\\simeq6$, i.e. the true energy of the fireball can be one order of magnitude larger than what inferred with the models of F01 and PK01. In addition (\\S~5) we can derive the GRBs luminosity function from the probability distribution of the viewing angle and compare it to data. The comparison is, at this stage, still uncertain because more accurate spectral and fluence measures are required to build a volume limited sample and confirm (or rule out) this model. In any fireball model considering a jet-like GRB structure, a certain number of afterglows without $\\gamma$-ray emission (orphan afterglows) is expected. For an homogeneous jet, orphan afterglows are possible only for viewing angles greater than $\\theta_j+1/\\Gamma$. In our jet configuration a fraction of the total area could have a Lorentz factor lower than the minimum $\\Gamma$ necessary for $\\gamma$-ray radiation. Consequently, considering the same opening angle, an inhomogeneous jet could produce an higher fraction of orphan afterglows than an homogeneous one. This intrinsic fraction depends above all on the $\\Gamma$ distribution within the jet ($\\alpha_{\\Gamma}$ in Eq~\\ref{eq:gammadis}) and on the minimum Lorentz factors to produce $\\gamma$-ray radiation and afterglow emission. The observed number of orphan afterglows depends also on flux detection limits, GRB explosion rates with redshift and cosmology. An accurate calculation of the expected orphan afterglow rates is therefore beyond the scope of this paper. We emphasize that we implicitly assumed the radiation efficiency of the fireball to be weakly dependent on $\\Gamma$ or $\\epsilon$. This is a plausible assumption in internal-shocks scenario. If it was not the case and the efficiency grew with $\\epsilon$, a different relation between $\\epsilon$ and $\\theta$ should be postulated in order to reproduced observations. \\begin{figure} \\psfig{file=MB1030_f4.ps,width=0.48\\textwidth} \\caption{{The lightcurves of inhomogeneous jets with different indices $\\alpha_\\epsilon$ ($\\epsilon\\propto\\theta^{-\\alpha_\\epsilon})$. The lightcurves have the same $\\epsilon_o$ but they are plotted shifted by factors of 10 for clarity (they would be indistinguishable at small time). The arrows highlight the location of the two breaks for $0<\\alpha<2$. } \\label{fig:alfae}} \\psfig{file=MB1030_f3.ps,width=0.48\\textwidth} \\caption{{The relation between the indices $\\alpha_t$ ($E_{iso}\\propto t_b^{-\\alpha_t}$) and $\\alpha_\\epsilon$ ($\\epsilon\\propto\\theta^{-\\alpha_\\epsilon}$) from an inhomogeneous jet. The shaded regions show the best fit to F01 data $\\alpha_t=0.9\\pm0.22$.} \\label{fig:indi}} \\end{figure} In this paper we concentrated for simplicity on a beam profile $\\epsilon\\propto\\theta^{-2}$ which is consistent observational results but it is interesting to briefly discuss other power-law relations $\\epsilon\\propto\\theta^{-\\alpha_\\epsilon}$ (see Fig.~\\ref{fig:alfae}). A decay flatter then 2 would cause two breaks in the light curve: the first due to the cone pointing the observer when $\\Gamma(\\theta_{o})\\sim\\theta_{o}^{-1}$ and the second, at later times, when the observer sees the edge of the jet and $\\Gamma(\\theta_{j}) \\sim \\theta_{j}^{-1}$. The power law index after the first break is flatter then $t^{-p}$ because the cones with $\\theta_{o}<\\theta\\leq\\theta_{j}$ enter the line of sight with $\\epsilon \\geq \\epsilon_{o}$ and substantially modify the light curve shape. With a steeper decay in the distribution of $\\epsilon$ the time break and the emission after that break would be dominated by the jet along the axis rather then by the very much weaker part directed to the observer. The jet break would then be preceded by a prominent flattening in the lightcurve, especially for $\\alpha_\\epsilon>3$, difficult to reconcile with observations. In this case we would have $\\gamma$-ray emission only for very small angles and the number of orphan afterglows would be much greater then expected from the $\\alpha_\\epsilon=2$ model. From Eq.~\\ref{eq:tb} we can derive the relation between the index $\\alpha_\\epsilon$ and $\\alpha_t$ (where $E_{iso}\\propto t_b^{\\alpha_t}$). We obtain $\\alpha_t=\\alpha_\\epsilon/2$ for $\\alpha_\\epsilon\\ge2$ and $\\alpha_t=3\\,\\alpha_\\epsilon/(8-\\alpha_\\epsilon)$ for $\\alpha_\\epsilon<2$, when the first break is considered. In Fig.~\\ref{fig:indi} we plot this relation overlaid on the interval in $\\alpha_t$ allowed by observations (we used F01 data). We derive $1.5\\lsim\\alpha_\\epsilon\\lsim2.2$, at the $1\\sigma$ level. Further $\\gamma$-ray and afterglow observations will allow to constrain this parameter much better in the future. Besides the luminosity function discussed in \\S~5, there are several ways in which this model can be proved or disproved. First, we have shown that the real total energy of the fireball can easily be an order of magnitude larger than what estimated by PK01 and F01. In this case, after the fireball has slowed down to mild-relativistic and sub-relativistic speed, radio calorimetry (Frail, Waxman \\& Kulkarni 2000) should allow us to detect the excess energy. In addition, in this model we naturally predict that the more luminous part of the fireball have higher Lorentz factors. This may help explaining the detected luminosity-variability (Fenimore \\& Ramirez-Ruiz 2000) and luminosity-lag (Norris, Marani \\& Bonnell 2000) correlations ( Salmonson 2000, Kobayashy, Ryde \\& MacFadyen 2001; Ramirez-Ruiz \\& Lloyd-Ronning 2002). Another constrain is given by polarization. Since fireball anisotropy is a basic ingredient of this model, inducing polarization (Ghisellini \\& Lazzati 1999; see also Sari 1999 and Gruzinov \\& Waxman 1999). The time evolution of the polarized fraction and of the position angle are however different from a uniform jet (Rossi et al. in preparation). Finally, the properties of the bursts should not depend on the location of the progenitor in the host galaxy, and therefore this model can accommodate the marginal detection of an $E_{\\rm iso}$-offset relation (Ramirez-Ruiz, Lazzati \\& Blain 2001) only if a distribution of $\\theta_j$ is considered." + }, + "0112/astro-ph0112426_arXiv.txt": { + "abstract": "A sample of $\\sim$ 20,000 galaxies covering 0.76 deg$^2$ were observed with the CFHT-UH8K up to $V<23.5$ and $I<22.5$. The angular correlation analysis of the red selected sample ($V-I>1.4$) shows a stronger amplitude than the blue selected sample at all cutoff magnitudes. This effect could be explained either by luminosity selection effects or by a true color segregation in which red objects are preferably found in denser regions. If the latter is true, and assuming that red galaxies are dominated by an older population of stars, this measurement supports the idea that galaxies evolve faster in the harsh environment of dense clusters than in the field, at fiducial redshifts of z$\\sim$0.5. ", + "introduction": " ", + "conclusions": "" + }, + "0112/astro-ph0112176_arXiv.txt": { + "abstract": "The masses and the evolutionary states of the progenitors of core-collapse supernovae are not well constrained by direct observations. Stellar evolution theory generally predicts that massive stars with initial masses less than about 30\\msol\\ should undergo core-collapse when they are cool M-type supergiants. However the only two detections of a SN progenitor before explosion are SN1987A and SN1993J, and neither of these was an M-type supergiant. Attempting to identify the progenitors of supernovae is a difficult task, as precisely predicting the time of explosion of a massive star is impossible for obvious reasons. There are several different types of supernovae which have different spectral and photometric evolution, and how exactly these are related to the evolutionary states of the progenitor stars is not currently known. I will describe a novel project which may allow the direct identification of core-collapse supernovae progenitors on pre-explosion images of resolved, nearby galaxies. This project is now possible with the excellent image archives maintained by several facilities and will be enhanced by the new initiatives to create Virtual Observatories, the earliest of which ({\\sc astrovirtel}) is already producing results. ", + "introduction": "Supernovae of Types\\,II and Ib/Ic are thought to occur during core collapse in massive stars at the end of their lifetimes. However the only definite detection of a SN progenitor is that of SN1987A in the LMC \\cite{white87}, which was a blue supergiant (B3I; \\opencite{wal89}). The progenitor of SN1993J in M81 was possibly identified as a K0\\,Ia star \\cite{alder94}. Neither progenitor is consistent with the canonical stellar evolution picture, where core-collapse occurs while the massive star is an M-supergiant. We still don't understand the physical mechanisms which underpin the different supernovae types, and how these are related to the evolution of the progenitor star. There is an understandable lack of observational data to constrain the last moments of stellar evolution. This conference is dedicated to understanding the basic building blocks of galaxy evolution, and supernova physics is a fundamental input parameter in determining the dynamical and chemical evolution of galaxies from the first stars in the Universe to present day gas-rich galaxies. Linking the observed supernova types to a stars initial mass, metallicity, binarity, environment and its subsequent evolution is not only important for those of us working on massive stellar evolution, it will also impact on galaxy evolution as a whole. ", + "conclusions": "" + }, + "0112/astro-ph0112340_arXiv.txt": { + "abstract": " ", + "introduction": "A primary goal for cosmology is to determine the total energy density of the Universe, $\\Omega$ and its constituents. It has been long recognised that this goal is achievable through the study of the redshift-distance relation of Type Ia supernovae (SNe). The relation is sensitive to different values of the cosmological parameters and this difference is more prominent at high redshifts. Therefore, the recently discovered SN at $z\\sim 1.7$ \\cite{riess2001cgunn} might prove to be invaluable in this respect. However, different systematic effects such as obscuration by grey dust, luminosity evolution of Type Ia SNe and gravitational lensing are also possibly more severe at high redshifts. In Ref.~\\cite{pointcgunn,rahmancgunn}, the systematic effects of gravitational lensing on a large sample of SNe have been investigated. Here, (with more details in Ref.~\\cite{uscgunn}), we investigate the effects of gravitational lensing due to galaxies lying close to the line-of-sight to SN1997ff, generalising the work in Refs.~\\cite{lewiscgunn} and \\cite{riess2001cgunn} by investigating the combined effects from a larger number of galaxies and estimating the masses and velocity dispersions of the lensing galaxies from the measured luminosities. Riess et al.~\\cite{riess2001cgunn} argued that the observed brightness of SN1997ff suggests that there cannot be a sizeable luminosity evolution for Type Ia's nor significant extinction by dust. Our work shows that the possible lensing magnification effects are large enough that the data is also consistent with an intrinsically dimmer supernova, or with significant dust density along the line-of-sight. ", + "conclusions": "\\label{sec:conclusions} Our lensing analysis show that a large range of magnifications of SN1997ff is possible for reasonable values of the galaxy masses and velocity dispersions. The value of the magnification is very sensitive to details in the modelling of the matter distribution in the lensing galaxies. The apparent (lack of) ellipticity of the host galaxy can be shown not to put any strong constraints on the magnitude of the magnification effect. Thus we conclude, that in order to use the apparent magnitude of a single high redshift SN to infer the values of any cosmological quantities, or even to place meaningful limits on the possible dimming of Type Ia SNe by intergalactic grey dust or luminosity evolution, very careful modelling of the galaxies along the line-of-sight is needed in order to control the systematic effects from lensing. This work has been carried out together with Edvard M\\\"ortsell and Ariel Goobar. The author would like to thank Lars Bergstr\\\"om, Joakim Edsj\\\"o and Peter Nugent for helpful discussions and Tomas Dahl\\'en for providing galaxy spectral templates. I would also like to thank the Swedish Research Council for financial support." + }, + "0112/astro-ph0112389_arXiv.txt": { + "abstract": "We investigate the structure and evolution of a geometrically thin viscous Keplerian circumbinary (CB) disk, using detailed models of their radiative/convective vertical structure. We use a simplified description for the evolution of the cataclysmic binary and focus on cases where the circumbinary disk causes accelerated mass transfer ($\\gta 10^{-8} \\mpy$). The inner edge of the disk is assumed to be determined by the tidal truncation radius and the mass input rate into the disk is assumed to be a small fraction ($10^{-5}$--$0.01$) of the mass transfer rate. Under the action of the viscous stresses in the disk the matter drifts outward with the optically thick region extending to several AU. The inner part of the disk is cool with maximum effective temperatures $\\lta 3,000$ K while the outermost parts of the disk are $\\lta 30$ K and optically thin. We calculate the effects of thermal instability on a sufficiently massive CB disk. It leads to outbursts reminiscent of those in thermally unstable accretion disks, with the instability remaining confined to the inner regions of the CB disk. However, for most of the evolutionary sequences the surface densities required to trigger instability are not reached. The spectral energy distributions from circumbinary disks are calculated, and the prospects for the detection of such disks in the infrared and submm wavelength regions are discussed. ", + "introduction": "Over the last decade evidence has been accumulating suggesting the presence of gaseous matter beyond the orbit of the components in cataclysmic variable binary systems (CVs). Its existence has been inferred from P-Cygni profiles in nova-like variables, super soft sources, and dwarf novae in outburst (see, for example, Deufel et al. 1999). In addition, the observation of narrow widths of single-peaked emission lines in SW Sex stars (see Thorstensen et al. 1991; Hellier 2000) and low velocity spectral features in AM CVn which do not follow the orbital motion of the binary system (see Solheim \\& Sion 1994) provide further support for this interpretation. Such a gaseous component may take the form of an extended shell or a flattened distribution of matter surrounding the binary system. Assuming that some fraction of this matter remains bound to the system and condenses into the orbital plane, in analogy to the compression of a radiation-driven wind into the equatorial plane in Be-stars (Bjorkman \\& Cassinelli 1993, Bjorkman \\& Wood 1995), a circumbinary disk may form. The presence of such a CB disk, as proposed by Spruit \\& Taam (2001), can provide an attractive explanation for some of the discrepancies between observations and the current theory of evolution of CVs. The importance of a CB disk for the evolution of a CV lies in the angular momentum it can extract from the binary, leading to mass transfer rates exceeding those due to magnetic braking and/or gravitational radiation. The binary exerts gravitational torques on the external disk, thereby transferring some of its orbital angular momentum. If some fraction of the mass lost from the main sequence-like secondary of the system forms a CB disk, a feedback process can operate to accelerate the binary evolution. As shown in the detailed computations by Taam \\& Spruit (2001) the evolution of the system can be significantly accelerated and the mass transfer rate from the secondary towards the white dwarf, $\\dot{M}_{\\rm 2}$, elevated to rates as high as $10^{-7} \\mpy$. In previous studies, we have focused on the response of the mass losing component and on the evolution of the binary system resulting from the influence of the CB disk, using a one zone approximation for its vertical structure. Here, we report on the results of detailed numerical calculations for the vertical structure and radial evolution of such disks surrounding a model CV. In order to determine the observability of such disks, their spectral energy distributions are also presented. The assumptions of the model computations are described in the next section. In \\S3 we describe the evolution of a CB disk neglecting angular momentum losses from the binary system. Results with the binary evolution taken into account are shown in \\S4 and several spectral energy distributions for some illustrative systems are presented in \\S5. We summarize and conclude in the last section. ", + "conclusions": "The evolution of a CB disk has been investigated for cases where it significantly accelerates the binary evolution, such that the mass transfer rate onto the white dwarf can reach $\\mdot \\gta 10^{-8} \\mpy$, representative of nova-like variables and super soft sources. The column densities and temperatures in these disks are found to be similar to those in circumstellar (CS) disks surrounding young stellar objects. The column densities decrease with radius and reach maxima in the innermost regions of the disk with $\\Sigma_{max}$ ranging up to $4 \\times 10^4$~g~cm$^{-2}$. The midplane temperatures range from $\\sim 3500$ K at the inner edge to less than about 30 K in the optically thin outer regions. In contrast to the CS disks, where accretion can take place steadily, the outflowing mass in CB disks is distinctly nonsteady. The outer boundary of the CB disk, $R_d$, increases with time (roughly as $R_d \\sim t^{1/2}$). We find that CB disks around CVs should extend to several AU, about 100 times smaller than CS disks. The numerical results show that the fractional mass input rate into the CB disk required to accelerate the evolution is lower than found in Taam \\& Spruit (2001). This is due to the effect of convection on the disk structure (ignored in Taam \\& Spruit), which causes the lower stable branch of the $\\nu-\\Sigma$ relation to extend to much higher column densities. We find that the CB disks can become thermally unstable, but only under rather special conditions, namely at high mass transfer rates combined with short orbital periods. Such conditions could be met in the short period double degenerate AM CVn systems or, in a different context, in disks around Be stars (Lee et al. 1991). For normal CVs, we find that the CB disk is unlikely to become unstable because of the long timescale needed to reach the critical column density. However, depending upon the response of the secondary to mass loss, the CB disk could become unstable during the accelerated evolution. If the CB disk is unstable, a heat front develops and propagates outward from the inner radius. In contrast to the accretion disks in dwarf novae, the heating front is quenched quickly because of the difficulty in raising the lower column densities encountered at larger radii to the increasing critical column densities needed to sustain propagation. The inner region cycles between the hot and cold states on a timescale of tens of years, inducing changes in the bolometric luminosity of a factor $\\sim$ 10. We find that the spectral energy distributions expected of CB disks in CVs can dominate the emission from the donor star and the accretion disk of the white dwarf at wavelengths $\\lambda \\gta 3 \\mu$m. At longer wavelengths the relative contribution from the CB disk to the total emission from the system increases. Typical flux levels for a nova-like variable for an assumed distance of 100 pc would be about 0.3 Jy at 10 $\\mu$m and a factor of 10 lower at $450 \\mu$m. Although the distances to CV's are not well known, the model calculations suggest that even in nova-like variables with distances order of 1 kpc, detection of CBs in the 10 - 20 $\\mu$m region should be possible using the infrared detectors on the Gemini Observatory (Telesco et al. 2001) and the Keck Observatory (Jones et al. 1998). Although the emission from these CB disks becomes optically thin at 350 $\\mu$m, it is possible that the longer wavelength radiation at $450 \\mu$m may be detectable using the SCUBA instrument. In our interpretation the nova-like variables would be CVs with mass transfer accelerated by a CB disk, which makes them natural targets for searches for CB disks. However, the predicted emission from the CB disk before the accelerated phase is similar, for a substantial period of time, to that during the accelerated phase. The possibility for their detection is therefore not limited to the systems with high mass transfer rates. We speculate that the `extra infrared components' discovered by Harrison et al. (2000) in several well studied dwarf novae may in fact be due to CB disks. These observations show that searches of CB disks surrounding bright cataclysmic variable systems are technically possible with current instrumentation." + }, + "0112/astro-ph0112206_arXiv.txt": { + "abstract": "Ongoing accretion of low-metallicity gas onto the disk is a natural prediction of semi-analytical Galactic chemical evolution models. This star formation fuel ameliorates the overproduction of metal-poor G- and K-dwarfs in the solar neighbourhood which otherwise plague so-called ``closed-box'' models of Galaxy evolution. Do High-Velocity Clouds (HVCs) represent the source of this necessary fuel? We know that HVCs provide an important clue as to the processes governing galaxy formation and evolution - what is less clear is whether their role lies more closely aligned with cosmology (as relics of the Local Group's formation) or star formation (as tidal debris from nearby disrupted dwarfs, or the waste byproducts of disk supernova-driven winds). I provide a summary of recent speculations as to the origins of HVCs, and highlight several future projects which will lead to a deeper understanding of the role they play in galaxy evolution. ", + "introduction": "Analytical models of Galactic chemical evolution invariably demand the existence of infalling (near?) pristine gas onto the disk in order to avoid the overproduction of low-metallicity stars - the so-called G-dwarf problem (Flynn \\& Morell 1997). Larson (1972) first suggested that this requisite star formation fuel might be associated with the population of High-Velocity Clouds (HVCs) seen moving at anomalous velocities with respect to differential Galactic rotation (Wakker \\& van Woerden 1997). Modern state-of-the-art chemical evolution models retain the need for this infalling fuel (e.g. Chiappini et~al. 1997), although most do not necessarily target the HVCs as the most likely culprit. Tosi (1988) suggests that infalling fuel more metal-rich than $\\ga$0.2\\,Z$_\\odot$ violates the present-day disk abundance constraints provided by HII regions, an hypothesis we are exploring with updated dual-infall (halo+disk phases) models (Chiappini et~al. 2002). Independent of these Galactic fuel arguments, a natural byproduct of hierarchical clustering galaxy formation scenarios (such as the currently favoured $\\Lambda$CDM) is that the halo of our Milky Way should be populated with $\\sim$500 satellites (Klypin et~al. 1999), and accretion of gas should continue (at some level) to the present-day (as in the aforementioned analytical models). This prediction is more than an order of magnitude discrepant with that actually observed ($\\sim$30 satellites). Blitz et~al. (1999) and Braun \\& Burton (1999) have both recently revived the classic hypothesis due to Verschuur (1969), suggesting that HVCs are Local Group interlopers. This Local Group Infall scenario is based upon the assumption that the gas we see as an HVC traces an underlying (dominant) dark matter halo. Do HVCs represent the reservoir of star formation fuel predicted to exist by analytical and numerical simulations of galaxy formation? Are they cosmological relics, waste byproducts related to tidal disruption of neighbouring dwarf galaxies, or supernova-driven ejecta from the disk? Whether they be cosmological, or related to star formation processes, HVCs represent crucial, yet mysterious, ingredients of galaxy formation. Under the former, the HI component of HVCs would contribute on the order of 10$^{11}$\\,M$_\\odot$ to the Local Group, while the latter\\footnote{Assuming typical distances of 4\\,kpc, and the integrated HI flux density from Table~3 of Wakker \\& van~Woerden 1991.} would correspond to order 10$^{7}$\\,M$_\\odot$ of HI in the halo. In what follows, I present a summary of the present-day state-of-affairs in HVC research, highlighting several intriguing (if sometimes confusing and/or contradictory) pieces of the puzzle, and avenues of future research. ", + "conclusions": "" + }, + "0112/astro-ph0112120_arXiv.txt": { + "abstract": "This paper reviews coagulation models for planet formation in the Kuiper Belt, emphasizing links to recent observations of our and other solar systems. At heliocentric distances of 35--50 AU, single annulus and multiannulus planetesimal accretion calculations produce several 1000 km or larger planets and many 50--500 km objects on timescales of 10--30 Myr in a Minimum Mass Solar Nebula. Planets form more rapidly in more massive nebulae. All models yield two power law cumulative size distributions, $N_C \\propto r^{-q}$ with $q$ = 3.0--3.5 for radii $r \\gtrsim$ 10 km and $N_C \\propto r^{-2.5}$ for radii $r \\lesssim$ 1 km. These size distributions are consistent with observations of Kuiper Belt objects acquired during the past decade. Once large objects form at 35--50 AU, gravitational stirring leads to a collisional cascade where 0.1--10 km objects are ground to dust. The collisional cascade removes 80\\% to 90\\% of the initial mass in the nebula in $\\sim$ 1 Gyr. This dust production rate is comparable to rates inferred for $\\alpha$ Lyr, $\\beta$ Pic, and other extrasolar debris disk systems. ", + "introduction": "Recent observations indicate that nearly all low and intermediate mass stars are born with massive circumstellar disks of gas and dust. Most young pre-main sequence stars with ages of $\\sim$ 1 Myr have gaseous disks with sizes of 100 AU or larger and masses of $\\sim$ 0.01 $M_{\\odot}$ \\citep{bec99,lad99}. Many older main sequence stars have dusty debris disks with sizes of 100--1000 AU \\citep{aum84,smi84,gai99,hab99,so00,spa01}. Current source statistics suggest the percentage of stars with observable disks declines from $\\sim$ 100\\% among the youngest stars to less than 10\\% for stars more than 1 Gyr old \\citep{bac93,art97,lad99,lag00}. Models for the formation of our solar system naturally begin with a disk. In the 1700's, Immanuel Kant and the Marquis de Laplace proposed that the solar system collapsed from a gaseous medium of roughly uniform density \\citep{kan55,lap96}. A flattened gaseous disk -- the protosolar nebula -- formed out of this cloud. The Sun contracted out of material at the center of the disk; the planets condensed in the outer portions. Although other ideas have been studied since Laplace's time, this picture has gained widespread acceptance. Measurements of the composition of the earth, moon, and meteorites support a common origin for the sun and planets \\citep[e.g.,][]{har76,and89}. Simulations of planet formation in a disk produce objects resembling known planets on timescales similar to the estimated lifetime of the protosolar nebula \\citep{saf69,gre78,wet93,pol96,ale98,lev98,kok00,kor00,cha01}. The Kuiper Belt provides a stern test of planet formation models. In the past decade, observations have revealed several hundred objects with radii of 50--500 km in the ecliptic plane at distances of $\\sim$ 35--50 AU from the Sun \\citep{jew93,luu96,gla97,jew98,chi99,luu00,gla01}. The total mass in these KBOs, $\\sim$ 0.1 $M_{\\oplus}$, suggests a reservoir of material left over from the formation of our solar system \\citep{edg49,kui51}. However, this mass is insufficient to allow the formation of 500 km or larger KBOs on timescales of $\\sim$ 5 Gyr \\citep{fer81,ste95,st97a,kl98}. The Kuiper Belt also provides an interesting link between local studies of planet formation and observations of disks and planets surrounding other nearby stars. With an outer radius of at least 150 AU, the mass and size of the Kuiper Belt is comparable to the masses and sizes of many extrasolar debris disks \\citep{bac93,art97,lag00}. Studying planet formation processes in the Kuiper Belt thus can yield a better understanding of evolutionary processes in other debris disk systems. Making progress on planet formation in the Kuiper Belt and the dusty disks surrounding other stars requires plausible theories which make robust and testable predictions. This paper reviews the coagulation theory for planet formation in the outer solar system \\citep[for reviews on other aspects of planet formation, see][]{man00}. After a short summary of current models for planet formation, I consider recent numerical calculations of planet formation in the Kuiper Belt and describe observational tests of these models. I conclude with a discussion of future prospects for the calculations along with suggestions for observational tests of different models of planet formation. ", + "conclusions": "The discovery of the Kuiper Belt in the 1990's provides fundamental constraints on models for the formation and evolution of planets in the outer parts of our solar system. The observations imply $\\sim 10^5$ KBOs with radii of 50--500 km and a total mass of $\\sim$ 0.1--0.2 $M_{\\oplus}$ beyond the orbit of Neptune. The theoretical challenge is to understand the formation of large objects in a current reservoir of material that is $\\sim$ 1\\% of the initial mass in the solar nebula. This goal assumes that KBOs formed locally and that the initial surface density of the solar nebula did not decease abruptly beyond the orbit of Neptune. Observations indicate typical disk radii of at least 100--200 AU in nearby pre-main sequence stars, which suggests that the disk of our solar system originally continued smoothly beyond the orbit of Neptune. Testing the assumption of local KBO formation relies on future comparisons between observations and theory. Coagulation calculations appear to meet the challenge posed by KBOs. Published numerical calculations demonstrate that the formation of a few Plutos and numerous 100--500 km KBOs in the outer parts of a solar system is inevitable \\citep[Stern 1995, 1996a;][]{st97a,dav99,kl99a}. For a variety of initial conditions, collisions between small bodies at 30--50 AU naturally produce larger objects. Once there is a range in sizes, dynamical friction efficiently reduces the orbital eccentricities of the largest objects. Large objects in nearly circular orbits grow quickly. At 30--50 AU, runaway growth can produce 100 km and larger objects on short timescales. These objects then grow slowly to radii of 1000 km or more. The initial disk mass sets the timescale for Pluto formation in the outer parts of a solar system. Objects grow faster in more massive disks. For single annulus calculations of planetesimals orbiting the Sun, the timescale to produce the first Pluto is \\begin{equation} t_P \\approx 20 ~ {\\rm Myr} ~ \\left ( \\frac{\\Sigma_{35}}{0.2 ~ {\\rm g ~ cm^{-2}}} \\right )^{-1} ~ , \\end{equation} where $\\Sigma_{35} \\approx 0.2 ~ {\\rm g ~ cm^{-2}}$ is the initial surface density of a minimum mass solar nebula model extrapolated into the Kuiper Belt at $\\sim$ 35 AU \\citep[Figure 2; see also][]{st97a,kl99a}. This timescale depends weakly on the initial conditions. Growth is more rapid in a solar nebula with small initial eccentricities and with small initial bodies \\citep{kl99a}. The growth timescale in the Kuiper Belt is smaller than expected from coagulation calculations in the inner solar system. \\citet{lis96} estimate a timescale to produce Moon-sized ($10^{26}$ g) objects as \\begin{equation} t_M \\approx {\\rm 0.5 ~ Myr} \\left ( \\frac{{\\rm 1~g~cm^{-2}}}{\\Sigma(a)} \\right ) \\left ( \\frac{a}{{\\rm 1 ~ AU}} \\right ) ^{3/2} ~ . \\end{equation} \\noindent This relation implies timescales of $\\sim$ 500 Myr at 35 AU and $\\sim$ 1 Gyr at 45 AU. Our single annulus models yield $t_M \\sim$ 100 Myr at 35 AU and $t_M \\sim$ 600 Myr at 70 AU. For calculations where the initial size distribution is composed of 1--10 km bodies, multiannulus models imply $t_M \\sim$ 200--300 Myr at 40--50 AU. Collisional damping causes the difference between our results and equation (13). In our calculations, collisional damping between small objects with radii of 1 m to 1 km reduces eccentricities by factors of 5--10. Dynamical friction couples the eccentricity reduction of the small bodies to the largest bodies. Because runaway growth begins when gravitational focusing factors are large, collisional damping in our Kuiper Belt models leads to an early onset of runaway growth relative to models of the inner solar system where the collisional evolution of small bodies is not important. Once large objects form in the outer part of a solar system, they stir up the velocities of small objects with radii of 10--100 km or less. Velocity stirring retards growth and produces debris. When the collision energy of small bodies is comparable to their tensile strength, the small bodies undergo a collisional cascade where planetesimals are ground down into smaller and smaller objects. This process produces numerous small grains which are ejected by radiation pressure ($\\lesssim$ 1--3 $\\mu$m grains) or pulled towards the Sun by Poynting-Robertson drag ($\\gtrsim$ 1--3 $\\mu$m grains). These grains are lost on short timescales of 1 Myr or less. When the collisional cascade begins, most of the mass in the outer solar system is contained in small objects that are easy to fragment. The collisional cascade thus robs the larger bodies of material. Because collisional cascades start sooner in the evolution when bodies are weaker, the size of the largest object in a calculation depends on the tensile strength of the small planetesimals. Our models yield Earth-sized objects in the Kuiper Belt for $S_0$ = $2 \\times 10^6$ erg g$^{-1}$ and Pluto-sized objects for $S_0$ = $10^2$ to $10^3$ erg g$^{-1}$. The theoretical models thus resolve the dilemma of large objects in a low mass Kuiper Belt. Runaway growth of small objects at 40--50 AU in the solar nebula places $\\sim$ 5\\%--10\\% of the initial mass in large objects with radii of 50--500 km or larger. The collisional cascade converts 80\\%--90\\% of the initial mass into debris which is removed from the Kuiper Belt on short timescales. Over the 4.5 Gyr lifetime of the solar system, gravitational interactions between KBOs and the gas giant planets can remove $\\sim$ 50\\% to 80\\% of the remaining mass. Given the uncertainties, collisions and dynamics appear capable of removing more than 90\\% of the original mass in the Kuiper Belt. The observed size distribution of KBOs provides strong observational tests of this picture. The final size distribution of a Kuiper Belt calculation has three components. The merger component at large sizes is a power-law with $q_f \\approx$ 3.0--3.5; the debris component at small sizes is a power law with $q_f \\simless$ 2.5. The collisional cascade depletes objects with intermediate sizes of 0.1--10 km. Depletion produces a dip in the size distribution for $S_0 \\lesssim 10^5$ erg g$^{-1}$. The observations of large KBOs generally agree with the power law slope predicted for the merger component. The data are consistent with $q$ = 3.3--3.5; the multiannulus models predict $q$ = 3.15--3.35. If dynamical interactions and collisional evolution continue to remove KBOs from the 40--50 AU annulus after 1 Gyr, the predicted number of KBOs is within a factor of two of the observed number of KBOs. The multiannulus calculations produce more KBOs with radii of 1000 km or larger than are observed with current surveys. The predicted number of these large objects depends on $S_0$ and is therefore uncertain. The observed number of large objects is plagued by small number statistics. Future surveys will provide robust constraints on the population of large objects. Improved multiannulus coagulation calculations which include dynamical interactions with gas giant planets will improve the predictions. Current constraints on the population of small KBOs are also consistent with model predictions. The data indicate a turnover in the KBO number counts, which implies a turnover in the size distribution for small objects. The derived turnover radius of 0.1--10 km is close to theoretical predictions. Better observations of the optical and far-IR surface brightnesses of the Kuiper Belt can provide better estimates of the slope of the size distribution for KBOs with radii of 1 mm to 1 m. Observations with larger telescopes may detect the turnover radius directly. Measuring the tensile strengths of comets provides an interesting test of this picture of KBO formation. In our models, the formation of Pluto by coagulation requires a tensile strength $S_0 \\gtrsim$ 400 erg g$^{-1}$. Large tensile strengths, $S_0 \\gtrsim$ $10^5$ erg g$^{-1}$, allow the formation of large bodies, $\\sim$ 2000--3000 km, which have not been detected in the outer solar system. Because objects with radii of 2000--3000 km can form before Neptune reaches its current mass, the lack of large KBOs implies $S_0 \\lesssim$ $10^4$ erg g$^{-1}$ in the coagulation theory. Estimates on the tensile strength derived from comet Shoemaker-Levy 9, $S_0 \\sim$ $10^2$ erg g$^{-1}$ to $10^4$ erg g$^{-1}$ \\citep[e.g.][]{gre95}, are close to the lower limit required to form Pluto. Theoretical estimates have a much larger range, $S_0 \\sim$ $10^2$ erg g$^{-1}$ to $10^6$ erg g$^{-1}$ \\citep{sir00}. As theoretical estimates improve and observations of disrupted comets become more numerous, these results can constrain the coagulation models. The coagulation calculations demonstrate that planet formation in the outer parts of other solar systems is also inevitable. The mass of a Minimum Mass Solar Nebula is comparable to the median disk mass derived for nearby pre-main sequence stars \\citep{bec99,lad99,man00}. The formation timescale for a 1000 km planet at 30--50 AU in one of these disks is therefore $\\sim$ 10--30 Myr. Although this planet cannot be observed directly, gravitational stirring leads to a collisional cascade and copious dust production. In the multiannulus models, dust is produced at a rate of roughly 0.1--1 Earth mass every 100 Myr \\citep[see also][]{kb02}. Observations of nearby debris disk systems are consistent with dust produced in a planet-forming disk. The sizes of debris disks, $\\sim$ 10--1000 AU, are similar to the radius of the Kuiper Belt. The ages of the youngest debris disk systems are comparable to the Pluto formation timescale of $\\sim$ 10--20 Myr \\citep{lag00}. If the timescale for Poynting-Robertson drag sets the residence time for 1 $\\mu$m and larger dust grains in the disk, the instantaneous dust mass in the disk is $\\sim$ 0.1--1 lunar masses. This mass is comparable to the dust masses inferred from IR observations of debris disk systems such as $\\alpha$ Lyr and $\\beta$ Pic \\citep{bac93,lag00}. Finally, the duration of the collisional cascade in our Kuiper Belt models, $\\sim$ 100 Myr to $\\sim$ 1 Gyr, is similar to the estimated lifetimes of debris disk systems, $\\sim$ 500 Myr \\citep[][2001]{hab99}. \\citet{kb01} derive a similar predicted lifetime for debris disk systems from the coagulation equation \\citep[see also][]{ken00}. To make the connection between KBOs and debris disks more clear, \\citet{ken99} investigate planet formation in the dusty ring of HR 4796A \\citep{jay98,koe98,au99,sch99,gr00}. They show that a planetesimal disk with a mass of 10--20 times the mass of the Minimum Mass Solar Nebula can form a dusty ring on 10--20 Myr timescales, comparable to the estimated age of HR 4796A. The model ring has a radial optical depth $\\sim$ 1, in agreement with limits derived from infrared images and from the excess infrared luminosity. Although the initial mass in this single annulus calculation is large, multiannulus calculations suggest similar timescales with much smaller masses. Finally, multiannulus calculations are an important new tool in developing a robust model for planet formation. Current computer technology allows practical multiannulus calculations that cover roughly a decade in disk radius. We are thus 1--2 orders of magnitude from constructing model grids of complete solar systems. Faster computers should resolve this difficulty in the next few years and allow us to consider the interfaces between (i) gas giants and terrestrial planets and (ii) gas giants and the Kuiper Belt. With some limitations, current multiannulus calculations promise predictions for the radial variation of the disk scale height \\citep{kb01} and the disk luminosity \\citep{kb02} as a function of stellar age, disk mass, and other physical parameters. Detailed comparisons between these predictions and observations of debris disks will yield interesting constraints on the physics of planet formation in other solar systems. Applying these results to our solar system will provide a better idea how the Earth and other planets in our solar system came to be. \\vskip 4ex I thank J. Luu for suggesting our joint projects and B. Bromley for advice and assistance in preparing the coagulation code for a modern, parallel computer. The JPL and Caltech supercomputer centers provided generous allotments of computer time through funding from the NASA Offices of Mission to Planet Earth, Aeronautics, and Space Science. Advice and comments from M. Geller, M. Kuchner, C. Lada, B. Marsden, R. Windhorst, and J. Wood greatly improved the content and the presentation of this review. \\vfill \\eject" + }, + "0112/astro-ph0112066_arXiv.txt": { + "abstract": "Analysis of the absorption lines in the afterglow spectrum of the gamma-ray burst GRB~010222 indicates that its host galaxy (at a redshift of z=1.476) is the strongest damped Lyman-$\\alpha$ (DLA) system known, having a very low metallicity and modest dust content. This conclusion is based on the detection of the red wing of Lyman-$\\alpha$ plus a comparison of the equivalent widths of ultraviolet Mg{\\smc I}, Mg{\\smc II}, and Fe{\\smc II} lines with those in other DLAs. The column density of H{\\smc I}, deduced from a fit to the wing of Lyman-$\\alpha$, is (5 $\\pm$ 2)\\ET{22} \\cm2. The ratio of the column densities of Zn and Cr lines suggests that the dust content in our line of sight through the galaxy is low. This could be due to either dust destruction by the ultraviolet emission of the afterglow or to an initial dust composition different to that of the diffuse interstellar material, or a combination of both. ", + "introduction": "Gamma Ray Bursts (GRBs), the most powerful cosmic explosions, are very important as probes of the high-redshift universe. This is due to several factors: (i) GRBs occur at cosmological distances, the most distant so far being at z=4.5 (GRB~000131, Andersen et al. 2000), (ii) they can be extremely bright and therefore could be seen (though for a very short period of time) up to redshift 20 (e.g. Wijers et al. 1998; Lamb \\& Reichart 2000), (iii) the burst and its afterglow illuminate the material situated between the GRB and the observer, providing an opportunity to study galaxies that otherwise would be too faint to be observed spectroscopically. This is particularly relevant for the host galaxies of GRBs where, in analogy to quasar (QSO) absorption line studies, the strengths of the absorption lines can be used to measure important properties, such as their gas content, dynamical structure, and metallicity. In addition, since they are transient phenomena, they do not cause the large scale ``proximity effect'' of quasars, i.e. they do not modify their environment at large distances, although they are expected to do so in their immediate vicinity (see section~\\ref{sec:discussion}), (iv) the currently most popular physical explanation of this phenomenon is the so-called fireball model (cf. Rees \\& Meszaros 1992) which is a relativistic explosion caused by either the collapse of a massive star (Woosley 1993; Paczynski 1998) or the merging of two compact objects (Lattimer \\& Schramm 1974). In the first case the rate of GRBs in the universe should be directly linked to the instantaneous star formation rate (SFR). This indeed seems to be supported by the observations, since the host galaxies of GRBs\\footnote{Actually, this is true for bursts whose duration is longer that 2 seconds, since they are the only ones which have been localized up to now.} are undergoing strong star-formation activity (e.g. Fruchter et al. 2001; Vreeswijk et al. 2001). Therefore, GRBs offer the potential to measure the cosmic star-formation rate (SFR) as a function of redshift. In this paper we report on spectroscopic observations of the GRB~010222 afterglow taken \\si 23 hours after the burst. We present evidence that its host galaxy is a damped Lyman-$\\alpha$ absorption (DLA) system with the strongest column density of neutral hydrogen and the lowest metallicity currently known. This has strong implications for both theories of the nature of GRBs and their host galaxies, and for the studies of DLA absorption systems. In particular, it shows that DLA systems with N(H{\\smc I}) $> 10^{22}$ \\cm2 {\\it do} exist, and that the way to detect them may be by looking at GRB afterglow spectra rather than QSO absorption line spectra. This is of extreme importance because DLAs are key components of the universe at high redshift due to the fact that, although relatively rare, they account for most of the neutral gas available for star formation. In section~\\ref{sec:obs} we present the observations and describe the calibration procedures. In section~\\ref{sec:analysis} the reduced spectra are analysed, resulting in a detailed identification of the absorption lines, and the determination of the metallicity and dust content of the gas producing them. Finally, in section~\\ref{sec:nature} we determine the nature of the galaxy hosting GRB~010222 and compare it to the host galaxies of other GRBs, as well as with other Mg{\\smc II} and damped Ly$\\alpha$ absorption systems. The discussion and conclusions are presented in section~\\ref{sec:discussion} and section~\\ref{sec:conclusions}, respectively. \\subsection{GRB~010222} GRB~010222 was discovered on February 22, 2001, at 7:23:30 U.T. by the Wide Field Camera (WFC) on board {\\it Beppo}SAX (Piro 2001). This burst is one of the brightest bursts ever recorded, with a fluence of (9.3 $\\pm$ 0.3)\\ET{-5} in the energy range from 40 to 700~keV. Soon after its detection, an optical transient was discovered with coordinates RA = 14\\hours 52\\minutes 12\\fseconds55 and DEC $=$+43\\arcdeg 01\\arcmin 06\\farcs2 (J2000), and magnitudes V\\si 18 and R $=$ 18.4\\mpm 0.1 (Henden 2001; McDowell et al. 2001; Henden \\& Vrba 2001). A series of optical spectra of the afterglow taken shortly after the alert allowed the determination of the redshift of the host galaxy (Garnavich et al. 2001; Jha et al. 2001; Bloom et al. 2001; Castro et al., 2001), which is $z=1.476$. For a cosmology with {H$_0=65$ \\kmsmpc, $\\Omega_m=0.3$, $\\Lambda=0.7$}, the luminosity distance is D$_L = $3.6\\ET{28} cm, and the isotropic energy output is 7.8\\ET{53} erg (in 't Zand et al. 2001), which makes GRB~010222 the third most energetic of all GRBs with known redshift. With regard to the host galaxy, observations of GRB~010222 at mm- and sub-mm-wavelengths have led to the interpretation that it is vigorously forming stars at a rate of $\\sim$500~M$_{\\odot}$~yr$^{-1}$ (Frail et al. 2001), whilst in the optical and near-infrared it is a faint (V$=$26.0\\mpm 0.1) compact (FWHM\\si 0\\farcs15) galaxy (Fruchter et al. 2001). ", + "conclusions": "\\label{sec:conclusions} The observational evidence points out that GRB~010222 has taken place in a damped Lyman-$\\alpha$ system having the highest column density of neutral hydrogen and the lowest metallicity known up to now, and with low dust content. This result reinforces the idea that a small cross-section, and not dust, is the cause of not finding DLA systems having very large N(H{\\smc I}). If the system associated with GRB~010222 is not atypical, this means that we may be missing an important fraction of H{\\smc I} at cosmological distances which in turn may have an impact on $\\Omega_{DLA}$ (i.e. the fraction of neutral gas in DLAs as a function of the closure density), although the exact contribution will depend on their cross section, which is unarguably small. It is not clear whether the dust is destroyed by the GRB itself, or if its size distribution is different, being devoided of small particles. In the latter case the feature due to silicate grains at 9.7 $\\mu$m should be absent also. Finally, there is strong evidence that other GRB host galaxies (or at least their locations) are as well damped Lyman-$\\alpha$ systems. Since they seem as well associated with star formation, we conclude that they may occur in a sub-class of DLAs that {\\it is} associated with active star formation." + }, + "0112/astro-ph0112250_arXiv.txt": { + "abstract": "We present a series of monitoring observations of the Crab Nebula with the {\\it Chandra X-ray Observatory}, focusing on the temporal evolution of the structure. This series of 8 observations, spanning a period of approximately six months, shows the dynamic nature of the inner X-ray structures. We detected outward moving ``wisps'' from the recently discovered inner ring seen in optical observations. We also find that the inner ring itself shows temporal variations in structure. The torus also appears to be expanding. Such temporal variations generally match the canonical scenario that an expanding synchrotron nebula injected from the pulsar is confined by the supernova ejecta. ", + "introduction": "\\vspace{-0.2cm} The Crab Nebula has been the best laboratory for investigating the mechanism linking the pulsar wind nebula (PWN) with the pulsar. The energetics confirms that the PWN is powered by the spin-down energy of the pulsar, and it is generally believed that this energy is transported by a relativistic wind (e.g. Kennel \\& Coroniti 1984). Hester (1998) and Tanvir, Thomson, \\& Tsikarishvili (1997) showed that the ``wisps'', which are elliptical ripples around the pulsar (Scargle 1969), are moving outwards with a speed of about 0.5{\\it c}. Such high energy phenomena must be associated with X-ray emission. Here, we present the results of a series of monitoring X-ray observations of the Crab Nebula with {\\it Chandra}, whose spatial resolution of $0.^{\\!\\!\\prime \\prime}5$ is comparable to that of ground-based optical telescopes. We adopt a distance of 2 kpc to the Crab Nebula throughout this paper. \\vspace{-0.3cm} ", + "conclusions": "\\vspace{-0.2cm} Figure~\\ref{fig:2ndimage} shows one of eight images of the Crab Nebula. It clearly shows an axisymmetric structure about the polar jet with the torus and the inner ring resolved in an early {\\it Chandra} observation (Weisskopf et al. 2000). Fine fibrous structures are also resolved at periphery of such large-scale structures. They show clear correlation with a optical polarization measurement (Hickson \\& van den Bergh 1990), indicating that they trace local magnetic field structures. \\begin{figure}[t] \\begin{center} \\hspace{0cm} \\psfig{file=fig1.ps,width=5.5cm} \\caption{{\\it Chandra} ACIS-S image of the Crab Nebula observed on 2000 November 25 (2nd observation), displayed with a square root brightness scale to enhance the faint structures. Due to heavy pileup which results in on-board event rejection, there is a hole where no events are detected at the position of the pulsar. The narrow line through the pulsar is the so-called trailed image of the pulsars.} \\label{fig:2ndimage} \\end{center} \\vspace{-0.5cm} \\end{figure} Figure~\\ref{fig:2468image} shows a series of the observations. Two wisps moving outward are detected through all of eight observations. They can also be seen in simultaneous HST observations (Hester et al. 2002). Here, we denote them as ``wisp A'' and ``wisp B''. They appear to break off from the inner ring. With respect to the inner ring's elliptical shape, the shapes of the wisps look warped due to the time delay of light travel (Hester et al. 2002). We measured the speed of the wisps, assuming that they are moving in the equatorial plane which includes the inner ring, and taking the inclination of the plane (the inclination angle was derived assuming the inner ring is circular) and the time delays into account. In spite of the difference in the directions and the birth times, the speeds of the two wisps are almost same, $\\sim$0.43{\\it c}. Their similarity indicates the existence of the continuous isotropic pulsar wind in the equatorial plane. \\begin{figure}[t] \\begin{center} \\hspace{0cm} \\psfig{file=fig2.ps,width=10cm} \\caption{Expanded images of the 2nd, 4th, 6th, and 8th observations, spaced 6 weeks apart. They demonstrate that two wisps are breaking off from the inner ring and that the inner ring varies.} \\label{fig:2468image} \\end{center} \\vspace{-0.6cm} \\end{figure} The inner ring is also variable, but unlike the wisps, it preserves its overall ring-like shape and relative position with respect to the pulsar. The ring never forms a continuous loop, but appears intermittent and mostly consists of knot-like features. Among them, three knots, which lie along the southeast portion of the ring and are symmetric about the axis of the jet, gradually brighten by factor $\\sim$1.5 within our observational period of 6 months. However, we note that these knots were bright enough to be detected 1 year before our 1st observation (Weisskopf et al. 2000). Additionally, some blob-like features appear to move outward along the jet, with a speed of $\\sim$0.34{\\it c}. \\begin{figure}[t] \\begin{center} \\hspace{0cm} \\psfig{file=fig3.ps,width=10cm} \\caption{Difference images of the 2nd$-$3rd (left), and the 2nd$-$8th (right) observations. Features which are strong in the earlier observation appear as white in the images. Longer time scale variations than that of the wisps can be seen in the torus.} \\label{fig:diff} \\end{center} \\vspace{-0.6cm} \\end{figure} Temporal variations can also be seen in the torus. Figure~\\ref{fig:diff} shows the difference images of the 2nd$-$3rd and the 2nd$-$8th observations. Although the variation is not strongly pronounced in the 3-week difference image, except for the wisps, differences within the torus and along its boundary are quite substantial over a duration of 19 weeks. The torus seems to be expanding at 0.1--0.2{\\it c}. Due to the small displacement and the ambiguous boundary, relatively large uncertainties remain. The fact that the angular extent along the major axis measured with {\\it Chandra} agrees well with those measured with {\\it Einstein} (Brinkmann et al. 1985), {\\it ROSAT} (Hester et al. 1995; Greiveldinger \\& Aschenbach 1999), and even from lunar occultation 25 years ago (Aschenbach \\& Brinkmann 1975) suggests that the torus is stable on tens of arcsecond scales over decades, but varies on arcsecond scales on several months. The northwestern region and the end of the jet do not exhibit any strong variations. Pavlov et al. (2001) discovered that the Vela PWN also shows temporal variability. Thus, it is indicated that the dynamic behavior is a common feature of PWNe. These X-ray observational pictures generally match the canonical scenario that the PWN is confined by the supernova ejecta. The relativistic winds from the pulsar are continuously transported in the equatorial plane and accumulate at the torus, outside of which the optical filaments -- ejecta -- can be seen." + }, + "0112/astro-ph0112299_arXiv.txt": { + "abstract": "We report on 3--4 $\\mu$m slit spectroscopy of 13 Seyfert 2 nuclei. The 3.3 $\\mu$m polycyclic aromatic hydrocarbon (PAH) emission is used to estimate the magnitudes of compact nuclear starbursts (on scales less than a few 100 pc) and to resolve the controversy over their energetic importance in Seyfert 2 nuclei. For three selected Seyfert 2 nuclei that have been well studied in the UV, the magnitudes of the compact nuclear starbursts estimated from the 3.3 $\\mu$m PAH emission (with no extinction correction) are in satisfactory quantitative agreement with those based on the UV after extinction correction. This agreement indicates that the flux attenuation of compact nuclear starburst emission due to dust extinction is insignificant at 3--4 $\\mu$m, and thus allows us to use the observed 3.3 $\\mu$m PAH luminosity to estimate the magnitudes of the compact nuclear starbursts in Seyfert 2 nuclei. Based directly on our 3--4 $\\mu$m slit spectra, the following two main conclusions are drawn: (1) except in one case, the observed nuclear 3--4 $\\mu$m emission is dominated by AGN and not by starbursts, and (2) compact nuclear starbursts are detected in 6 out of 13 Seyfert 2 nuclei, but cannot dominate the energetics of the galactic infrared dust emission in the majority of the observed Seyfert 2 galaxies. For several sources for which Infrared Space Observatory spectra taken with larger apertures and/or soft X-ray data are available, these data are combined with our 3--4 $\\mu$m slit spectra, and it is suggested that (3) extended (kpc scale) star-formation activity is energetically more important than compact nuclear starbursts, and contributes significantly to the infrared luminosities of Seyfert 2 galaxies, although the AGN is still an important contributor to the luminosities, and (4) the bulk of the energetically significant extended star-formation activity is of starburst type rather than quiescent normal disk star-formation; the extended starbursts are responsible for the superwind-driven soft X-ray emission from Seyfert 2 galaxies. Finally, a correlation between the luminosities of AGNs and compact nuclear starbursts is implied; more powerful AGNs tend to be related to more powerful compact nuclear starbursts. ", + "introduction": "According to the unification paradigm for Seyfert galaxies, Seyfert 1s, which show broad optical emission lines, and 2s, which do not, are intrinsically the same objects, but the nuclei of the latter class are obscured by dust along our line of sight in dusty molecular tori \\citep{ant93}. It is widely accepted that the ultimate energy source of Seyfert {\\it nuclei} (although not of their extended host galaxies) is the release of gravitational energy caused by mass accretion onto a central supermassive blackhole (so-called AGN activity). However, it has recently been argued that strong signatures of starbursts are detected in the nuclear UV--optical spectra of Seyfert 2 galaxies (Heckman et al. 1997; Gonzalez Delgado et al. 1998; Storchi-Bergmann et al. 2000; Gonzalez Delgado, Heckman \\& Leitherer 2001), and that the intrinsic extinction-corrected luminosities of compact nuclear starbursts (hereafter `compact' is used to mean a size scale of less than a few 100 pc) could be comparable to the luminosities of the AGN \\citep{gon98}. In contrast, \\citet{iva00} investigated the CO indices in the near-infrared $K$-band spectra of Seyfert 2 nuclei, and found no evidence for the presence of strong compact nuclear starbursts. Although these authors observed different samples of Seyfert 2 galaxies, it is nevertheless striking that they draw such contradictory conclusions on the energetic role of compact nuclear starbursts in Seyfert 2 nuclei. According to \\citet{cid95}, compact nuclear starbursts are natural byproducts of AGNs' dusty molecular tori. In this case, since direct UV--optical emission from the AGN is highly attenuated in Seyfert 2 nuclei, the less strongly obscured compact nuclear starburst emission will inevitably make a relatively strong contribution to the observed UV--optical fluxes in these objects. Thus, the detection of signatures of compact nuclear starbursts in the observed UV--optical spectra of Seyfert 2 nuclei \\citep{sto00,gon98,gon01} would not be surprising. However, to obtain a deeper understanding of the nature of Seyfert 2 nuclei, it is certainly more important to quantify the energetic importance of compact nuclear starbursts than simply to investigate their presence. \\citet{gon98} estimated the magnitudes of compact nuclear starbursts, based on their extinction-corrected UV data, for four UV-bright Seyfert 2 galaxies, while for many other Seyfert 2 galaxies that have been studied optically, no quantitative discussions of the absolute magnitudes of compact nuclear starbursts have been made \\citep{sto00,gon01}. Even the UV-based quantitative discussion could contain large uncertainties, because of the susceptibility of UV emission to dust extinction; the extinction correction factor for UV data may vary drastically depending on the assumed amount of dust and its spatial distribution, and has often been claimed to be quantitatively uncertain in particular applications \\citep{bak01}. The uncertainties in the dust extinction correction are much less in the near-infrared $K$-band data taken by Ivanov et al. However, the spectral signatures of starbursts are weak, so that careful subtraction of old stellar and AGN emission is required to estimate the magnitudes of compact nuclear starbursts \\citep{iva00}. At 3--4 $\\mu$m, we possess powerful diagnostic tools to distinguish between starburst and AGN activity, to detect weak starbursts, and to estimate the magnitudes of compact nuclear starbursts in Seyfert 2 galaxies with fewer quantitative uncertainties. As shown by Imanishi \\& Dudley (2000, their figure 1), \\begin{itemize} \\item[(a)] If a Seyfert 2 nucleus is powered by starbursts, strong polycyclic aromatic hydrocarbon (PAH) emission will be detected at 3.3 $\\mu$m regardless of dust extinction. \\item[(b)] If it is powered by obscured AGN activity, a carbonaceous dust absorption feature will be detected at 3.4 $\\mu$m. \\item[(c)] If it is powered by weakly obscured AGN activity, the 3--4 $\\mu$m spectrum will be nearly featureless. \\item[(d)] If it is a composite of starburst and AGN activity, the absolute luminosity and the equivalent width of the 3.3 $\\mu$m PAH emission feature will be smaller than in starburst galaxies, so that these values can be used to quantitatively estimate the energetic importance of starburst activity. \\end{itemize} Since the 3.3 $\\mu$m PAH emission is intrinsically strong \\citep{moo86,imd00}, even the signatures of weak starbursts are detectable; even if only $\\sim$10 \\% of the observed nuclear 3--4 $\\mu$m flux originates in starbursts, the 3.3 $\\mu$m PAH emission peak is $\\sim$20 \\% higher than the continuum level, and thus is clearly recognizable in normal S/N $\\sim$ 15--20 spectra. Additionally, the effects of dust extinction are smaller at 3--4 $\\mu$m than at shorter wavelengths, which makes the uncertainties in the dust extinction correction factor much smaller. This advantage is demonstrated quantitatively below. Suppose that the dust extinction in starbursts is found to be A$_{\\rm V}$ = 2--4 mag, with an uncertainty of a factor of 2 in A$_{\\rm V}$. Since the dust extinction in the UV ($\\sim$2000 \\AA) is $\\sim$2.5 $\\times$ A$_{\\rm V}$ \\citep{sav79}, the corresponding dust extinction correction factor in the UV is 100--10000, differing by a factor of 100 depending on the adopted A$_{\\rm V}$. The difference becomes even larger if we go to shorter-wavelength parts of the UV region or if the actual dust extinction is larger. In contrast, since the dust extinction at 3--4 $\\mu$m is $\\sim$0.05 $\\times$ A$_{\\rm V}$ \\citep{rie85,lut96}, the extinction correction factor at 3--4 $\\mu$m in the case of A$_{\\rm V}$ = 2--4 mag is 1.1--1.2. Thus, the extinction correction is negligible at 3--4 $\\mu$m, and the uncertainty in the correction factor is only at the 10\\% level. The high detectability of weak starbursts and the small uncertainties in dust extinction correction at 3--4 $\\mu$m combine to make 3--4 $\\mu$m observations a very powerful tool to quantitatively address the issue of the energetics of compact nuclear starbursts in Seyfert 2 nuclei. This paper reports the results of 3--4 $\\mu$m spectroscopy of Seyfert 2 nuclei and their implications for the energetics of these objects. Throughout the paper, $H_{0}$ $=$ 75 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm M}$ = 0.3, and $\\Omega_{\\rm \\Lambda}$ = 0.7 are adopted. ", + "conclusions": "\\subsection{What Powers the Observed Compact Nuclear 3--4 $\\mu$m Emission?} The fraction of the observed nuclear 3--4 $\\mu$m fluxes in our slit spectra that originates in starbursts can be estimated from the rest-frame equivalent widths of the 3.3 $\\mu$m PAH emission feature (EW$_{3.3 \\rm PAH}$). The EW$_{3.3 \\rm PAH}$ value decreases as the AGN contribution increases ($\\S$ 1). The equivalent widths for starburst-dominated galaxies are $\\sim$120 nm \\citep{moo86,imd00}. In Table~\\ref{tbl-4}, only Mrk 266SW shows an EW$_{3.3 \\rm PAH}$ value similar to those of starburst galaxies; we can conclude from this that only the observed nuclear 3--4 $\\mu$m flux of Mrk 266SW is dominated by starbursts. Mrk 78, Mrk 273, Mrk 477, NGC 3227, and NGC 5135 all show detectable, moderately strong 3.3 $\\mu$m PAH emission features, but their EW$_{3.3 \\rm PAH}$ values are a factor 3--9 smaller than those of starburst galaxies. For these five sources, starbursts contribute some fraction, but less than half, of the observed nuclear 3--4 $\\mu$m fluxes, and the bulk of the observed 3--4 $\\mu$m fluxes originate in AGN activity. For the remaining seven Seyfert 2 nuclei (IC 3639, Mrk 34, Mrk 463, Mrk 686, NGC 1068, NGC 5033, and NGC 5929), the EW$_{3.3 \\rm PAH}$ is more than a factor of 7 smaller than that expected for starburst-dominated galaxies, indicating that the predominant fraction (larger than 80\\%) of their 3--4 $\\mu$m emission comes from AGN activity. {\\it Our first conclusion is that, except in the case of Mrk 266SW, the observed compact nuclear 3--4 $\\mu$m emission is dominated by the AGN and not by starbursts}. This result has implications for infrared studies of AGNs. $L$- (3.5 $\\mu$m) and $M$-band (4.8 $\\mu$m) emission from AGNs is usually dominated by compact emission, and photometric data at these bands are used to discuss dust obscuration toward AGNs, on the assumption that this compact emission comes predominantly from the AGN and not from compact nuclear starbursts \\citep{sim98,ima01}. Our results provide supporting evidence for this assumption. \\subsection{The 3.4 $\\mu$m Dust Absorption Feature} The 3.4 $\\mu$m carbonaceous dust absorption feature is detected in Mrk 463 and NGC 1068. The observed 3.4 $\\mu$m absorption optical depths are $\\tau_{3.4}$(observed) $=$ 0.042$\\pm$0.005 and 0.12$\\pm$0.01 for Mrk 463 and NGC 1068, respectively. In AGNs, the 3--4 $\\mu$m continuum emission is dominated by dust in thermal equilibrium with a temperature of 800--1000K, close to the dust sublimation temperature, located at the innermost part of the dusty torus \\citep{sim98,alo98}. The dust extinction toward the 3--4 $\\mu$m continuum emitting region is thus almost the same as that toward the AGN itself. Assuming a Galactic dust model ($\\tau_{3.4}$/A$_{\\rm V}$ $=$ 0.004--0.007; Pendleton et al. 1994), the dust extinction toward the AGNs is estimated to be 6--11 and 17--30 mag for Mrk 463 and NGC 1068, respectively. Mrk 463 shows a broad emission component in its near-infrared hydrogen recombination lines \\citep{good94,vei97} so that dust obscuration toward the AGN has been suggested to be relatively modest in this object. Our result is consistent with this suggestion. For the remaining 11 Seyfert 2 nuclei, no detectable 3.4 $\\mu$m absorption feature is found. The 3.4 $\\mu$m absorption feature may be suppressed if the bulk of the dust grains in the obscuring material is covered with an ice mantle \\citep{men01}. However, no strong 3.08 $\\mu$m ice absorption feature is found in their 3--4 $\\mu$m spectra. The dust absorption feature is smeared out if less obscured starburst emission contributes significantly to the observed nuclear 3--4 $\\mu$m flux \\citep{ima00}. The six sources with detectable PAH emission (Mrk 78, Mrk 266SW, Mrk 273, Mrk 477, NGC 3227, and NGC 5135) could be this case. For the five PAH-undetected Seyfert 2 nuclei (IC 3639, Mrk 34, Mrk 686, NGC 5033, and NGC 5929), the 3.4 $\\mu$m absorption feature could remain undetectable if dust obscuration toward these AGNs is weak ($\\S$ 1). The R-value, defined as \\begin{eqnarray} R \\equiv log \\frac{\\nu_{3.5 \\mu m} F(\\nu_{3.5 \\mu m})}{\\nu_{25 \\mu m} F(\\nu_{25 \\mu m})} \\end{eqnarray} is a good measure for dust extinction toward AGNs (Murayama, Mouri, \\& Taniguchi 2000). It becomes smaller with higher dust extinction toward AGNs, and the boundary value between Seyfert 1 and 2 nuclei is estimated to be R = $-$0.6 \\citep{mura00}. Based on the 3.5 $\\mu$m (Table \\ref{tbl-3}) and {\\it IRAS} 25 $\\mu$m (Table \\ref{tbl-1}) fluxes, the R-values for IC 3639, Mrk 34, Mrk 686, NGC 5033, and NGC 5929, are estimated to be $-$1.2, $-$0.8, $-$0.3, $-$0.6, and $-$1.2, respectively. The R-values of Mrk 686 and NGC 5033 are in the range of Seyfert 1 nuclei, and thus weak dust obscuration is indicated, explaining the non-detection of the 3.4 $\\mu$m absorption. However, IC 3639, Mrk 34, and NGC 5929 show small R-values. The non-detection could be caused by the effects of time variability or contamination from extended star-formation to the observed 25 $\\mu$m flux taken with large apertures \\citep{alo01}. \\subsection{The Magnitudes of Compact Nuclear Starbursts as Estimated from the 3.3 $\\mu$m PAH Emission Fluxes} The bulk of the UV--optical emission from the AGNs in Seyfert 2 galaxies and starbursts \\citep{cal00} is absorbed by dust and re-emitted as dust thermal emission in the infrared. Infrared (8--1000 $\\mu$m) luminosities are thus used to estimate the magnitudes of AGN and starburst activity throughout this paper. Infrared (8--1000 $\\mu$m) luminosities are obviously better for evaluating these kinds of activity than far-infrared (40--500 $\\mu$m) luminosities because not all activity shows dust emission peaking in the far-infrared. In starburst-dominated galaxies, the 3.3 $\\mu$m PAH to far-infrared (40--500 $\\mu$m) luminosity ratios (L$_{3.3 \\rm PAH}$/L$_{\\rm FIR}$) are $\\sim$1 $\\times$ 10$^{-3}$ \\citep{mouri90}. The 3.3 $\\mu$m PAH to infrared (8--1000 $\\mu$m) luminosity ratio (L$_{3.3 \\rm PAH}$/L$_{\\rm IR}$) is tentatively assumed to be also 1 $\\times$ 10$^{-3}$ for the compact nuclear starbursts in Seyfert 2 nuclei. The scatter of the PAH-to-infrared luminosity ratio is likely to be a factor of 2--3 toward both higher and lower values around a typical value \\citep{fis00}. For IC 3639, Mrk 477, and NGC 5135, \\citet{gon98} explicitly estimated the extinction-corrected luminosities of compact nuclear starbursts, based on their UV data. The respective luminosities are 1.1 $\\times$ 10$^{43}$, 1.4 $\\times$ 10$^{44}$, and 4.2 $\\times$ 10$^{43}$ ergs s$^{-1}$. The infrared luminosities of the compact nuclear starbursts estimated from the observed (extinction-uncorrected) 3.3 $\\mu$m PAH emission, based on the assumption of L$_{3.3 \\rm PAH}$/L$_{\\rm IR}$ $\\sim$ 1 $\\times$ 10$^{-3}$, are $<$1.2 $\\times$ 10$^{43}$, 1.5 $\\times$ 10$^{44}$, and 5.8 $\\times$ 10$^{43}$ ergs s$^{-1}$, respectively, for IC 3639, Mrk 477, and NGC 5135. It is noteworthy that the compact nuclear starburst luminosities estimated from UV data after extinction correction and from the 3.3 $\\mu$m PAH emission with {\\it no} extinction correction are in satisfactory agreement for all three sources. This agreement indicates that (1) discussions of the magnitudes of compact nuclear starbursts based on UV data after extinction correction are quantitatively reliable, (2) the extinction correction factor is not significant at 3--4 $\\mu$m, and (3) the assumption of L$_{3.3 \\rm PAH}$/L$_{\\rm IR}$ $\\sim$ 1 $\\times$ 10$^{-3}$ is reasonable. We have been concerned by the possibility that, if compact nuclear starbursts close to the central AGNs are directly exposed to X-ray emission from AGNs, PAH emission from the compact nuclear starbursts might be suppressed due to the destruction of PAHs \\citep{voit92}. However, the agreement suggests that this is not the case. If the bulk of the compact nuclear starbursts in Seyfert 2 nuclei occur at the outer reaches of the dusty tori around AGNs \\citep{hec97}, PAHs are shielded from the energetic radiation from the AGNs and thus are not destroyed, explaining the strong PAH emission from compact nuclear starbursts in Seyfert 2 nuclei. Based on these findings, the magnitudes of compact nuclear starbursts can be estimated from the observed 3.3 $\\mu$m PAH emission luminosities with some confidence. If compact nuclear starbursts energetically dominated the whole galactic infrared dust emission luminosities of Seyfert 2 galaxies, the observed L$_{3.3 \\rm PAH}$/L$_{\\rm IR}$ ratios measured from our slit spectra should be similar to those of starburst galaxies ($\\sim$1 $\\times$ 10$^{-3}$). As shown in Table~\\ref{tbl-4}, however, none of the observed Seyfert 2 galaxies show such a large value of L$_{3.3 \\rm PAH}$/L$_{\\rm IR}$. The differences are by a factor of 3 (Mrk 477) to larger than 10, and in most cases much larger than a factor of 2--3, the possible flux loss in our slit spectra and the intrinsic scatter of the L$_{3.3 \\rm PAH}$/L$_{\\rm IR}$ ratios for starbursts. Therefore, {\\it our second conclusion is that compact nuclear starbursts cannot be the dominant energy source of the whole galactic infrared dust emission luminosities in the majority of the observed Seyfert 2 galaxies}. \\subsection{Comparison with Optical and $K$-band Diagnostic Results for Individual Sources} For the remaining ten Seyfert 2 galaxies other than IC 3639, Mrk 477, and NGC 5135, our quantitative estimates of the magnitudes of the compact nuclear starbursts are compared with qualitative arguments based on the optical and $K$-band data \\citep{gon01,iva00}. \\citet{gon01} found signatures of compact nuclear starbursts (young stars) in the optical spectrum of Mrk 273. For Mrk 273, 3.3 $\\mu$m PAH emission is detected; EW$_{3.3 \\rm PAH}$ and L$_{3.3 \\rm PAH}$/L$_{\\rm IR}$ are a factor of $\\sim$3 and $\\sim$10 smaller, respectively, than those of starburst galaxies. The L$_{3.3 \\rm PAH}$/L$_{\\rm IR}$ ratio indicates that the compact nuclear starbursts explain $\\sim$1/10 of the total galactic infrared luminosity. Thus, compact nuclear starbursts are certainly present, supporting the optical results. For Mrk 78 and NGC 1068, \\citet{gon01} found no signatures of compact nuclear starbursts, and detected only intermediate age stars in their optical spectra. For NGC 1068, only stringent upper limits are found for the EW$_{3.3 \\rm PAH}$ and L$_{3.3 \\rm PAH}$/L$_{\\rm IR}$, both of which are more than an order of magnitude smaller than those of starburst galaxies. For Mrk 78, however, 3.3 $\\mu$m PAH emission is marginally detected. For Mrk 463, \\citet{gon01} claimed that compact nuclear starbursts might be present, subject to further confirmation. No detectable 3.3 $\\mu$m PAH emission is found. The upper limits for both EW$_{3.3 \\rm PAH}$ and L$_{3.3 \\rm PAH}$/L$_{\\rm IR}$ are more than an order of magnitude smaller than starburst galaxies. Compact nuclear starbursts in Mrk 463 are energetically insignificant, if they exist at all. For Mrk 34, \\citet{gon01} found no signs of compact nuclear starbursts in the optical spectrum, with only old population stars being detected. No 3.3 $\\mu$m PAH emission is detected either, supporting the conclusions based on the optical data. Both \\citet{gon01} and \\citet{iva00} studied NGC 5929 and found no evidence for strong compact nuclear starbursts. No clear 3.3 $\\mu$m PAH emission is detected. The remaining sources, Mrk 266SW, Mrk 686, NGC 3227, and NGC 5033, were studied only by \\citet{iva00}. They found that no strong compact nuclear starbursts are required in terms of the CO indices after correction for AGN emission. No 3.3 $\\mu$m PAH emission is detected in Mrk 686 and NGC 5033, supporting their argument. However, in Mrk 266SW, the nuclear 3--4 $\\mu$m emission must be dominated by starbursts, although these detected compact nuclear starbursts can explain only $\\sim$10\\% of the total galactic infrared luminosity. 3.3 $\\mu$m PAH emission was also clearly detected in NGC 3227, and compact nuclear starbursts can explain $\\sim$20\\% of the whole galactic infrared luminosity. In Mrk 266SW and NGC 3227, compact nuclear starbursts are certainly present, and they explain a non-negligible fraction of the total galactic infrared luminosities. The signatures of compact nuclear starbursts may have been missed during the complicated procedures required to find their weak signatures in the $K$-band \\citep{iva00}. \\subsection{The Energetic Importance of AGN and Extended Star-Formation Activity} Since the observed nuclear 3--4 $\\mu$m emission is dominated by AGN activity ($\\S$ 5.1), it would be expected that AGN activity also contributes much more than compact nuclear starbursts to the net galactic infrared (8--1000 $\\mu$m) luminosities of Seyfert 2 galaxies. However, the AGN-driven infrared luminosities of these objects are difficult to estimate because they are highly dependent on the spatial distribution of dust in the dusty torus, which is poorly constrained observationally. Thus, we estimate the energetic importance of extended star-formation in the host galaxies, which is the remaining probable power source of Seyfert 2 galaxies other than compact nuclear starbursts and AGNs, by using the PAH emission. PAH emission can be produced both with starbursts and normal quiescent star-formation in a similar way \\citep{hel00}. To investigate the magnitude of extended (kpc scale) star-formation activity, {\\it Infrared Space Observatory (ISO)} spectra taken with large apertures are useful. \\citet{gh00} summarized the {\\it ISO} results; they found that (1) the 7.7 $\\mu$m PAH to far-infrared luminosity ratios in Seyfert galaxies are similar to those of starburst galaxies, and that (2) the bulk of the PAH emission is spatially extended. Therefore, they argued that the bulk of the far-infrared emission from Seyfert galaxies originates in extended star-formation activity. Of the 13 Seyfert 2 galaxies, \\citet{cla00} have presented {\\it ISO} 2.5--11 $\\mu$m spectra taken with 24 $\\times$ 24 arcsec$^{2}$ apertures, and quote PAH fluxes for four sources (Mrk 266, NGC 3227, NGC 5033, and NGC 5929). For these sources, the energetic importance of extended star-formation activity can be investigated by comparing the measured PAH fluxes in our slit spectra with those in the {\\it ISO} spectra. The effects of dust extinction are similar at 3--8 $\\mu$m \\citep{lut96}. The 7.7 $\\mu$m PAH emission is the strongest PAH emission feature. However, its flux estimates in the {\\it ISO} spectra may be highly uncertain \\citep{cla00,lau00}, due to the presence of strong, spectrally broad 9.7 $\\mu$m silicate dust absorption feature and insufficient wavelength coverage longward of these emission and absorption features, which make a continuum determination difficult \\citep{dud99}. Quantitatively reliable flux estimates of the 3.3 $\\mu$m PAH emission in the {\\it ISO} spectra are also difficult due to the scatter in the data points at 3--4 $\\mu$m (Clavel et al. 2000, their figure 8). Consequently, the 6.2 $\\mu$m PAH emission, which is isolated and moderately strong, is most suitable to investigate the extended star-formation activity based on the {\\it ISO} spectra \\citep{fis00}. The 6.2 $\\mu$m PAH to infrared luminosity ratios (L$_{6.2 \\rm PAH}$/L$_{\\rm IR}$) for starbursts are estimated to be $\\sim$ 6 $\\times$ 10$^{-3}$, with a scatter of a factor of 2--3 toward both higher and lower values \\citep{fis00}, where it is assumed that L$_{\\rm IR}$ $\\sim$ L$_{\\rm FIR}$ for starbursts. The respective L$_{6.2 \\rm PAH}$/L$_{\\rm IR}$ ratios are 4 $\\times$ 10$^{-3}$, 4 $\\times$ 10$^{-3}$, 3 $\\times$ 10 $^{-3}$, and 1 $\\times$ 10$^{-3}$, for Mrk 266, NGC 3227, NGC 5033, and NGC 5929 \\citep{cla00}, roughly half of (Mrk 266, NGC 3227, and NGC 5033) or more than six times smaller than (NGC 5929) the typical value for systems dominated by starbursts. Since the observed L$_{3.3 \\rm PAH}$/L$_{\\rm IR}$ ratios indicate that the compact nuclear starbursts energetically fall short of the total infrared luminosities by a factor of larger than 6 for these four sources (Table~\\ref{tbl-4}), it can be said that extended star-formation (but not compact nuclear starbursts) contributes significantly to the infrared luminosities of these Seyfert 2 galaxies. This statement was made by \\citet{gh00} for Seyfert galaxies as a whole based on the 7.7 $\\mu$m PAH emission. Based on the 6.2 $\\mu$m PAH emission, here, it has been confirmed that this is true also for the four individual Seyfert 2 galaxies, and probably for Seyfert 2 galaxies as a whole. The individual L$_{6.2 \\rm PAH}$/L$_{\\rm IR}$ ratios for Mrk 266, NGC 3227, and NGC 5033 are within the scattered range of the ratios for starbursts. However, the overall trend of lower values implies that AGN activity also makes an important contribution to the infrared luminosities of Seyfert 2 galaxies. \\subsection{Is Extended Star-Formation Quiescent or of Starburst Type?} While both starbursts and quiescent disk star-formation in normal galaxies can produce PAH emission ($\\S$ 5.5), strong soft X-ray emission driven by superwind is observed only if the star-formation rate per unit area exceeds a certain threshold (10$^{-1}$ M$_{\\odot}$ yr$^{-1}$ kpc$^{-2}$; Heckman 2000). Starbursts surpass this threshold, while quiescent normal disk star-formation does not \\citep{ken98}. These different characteristics can be used to understand the properties of energetically significant extended star-formation activity. Of the 13 Seyfert 2 galaxies, \\citet{lev01} have estimated superwind-driven soft X-ray luminosities for eight sources (IC 3639, Mrk 78, Mrk 266, Mrk 273, Mrk 463, Mrk 477, NGC 1068, and NGC 5135), and found that, as a whole, their soft X-ray to far-infrared luminosity ratios are as high as those of starburst galaxies. Therefore, starburst activity is a significant contributor to the far-infrared emission and also to the infrared dust emission. Since compact nuclear starbursts detected in our slit spectra were found to be energetically insignificant ($\\S$ 5.3), the energetically significant starbursts must be extended. {\\it Our third conclusion is that the bulk of the energetically significant extended (kpc scale) star-formation activity is of starburst-type and not quiescent normal disk star-formation. It is the extended (kpc scale) starbursts, rather than the compact (less than a few 100 pc) nuclear starbursts, that are responsible for the superwind-driven soft X-ray emission.} If starbursts are inevitable phenomena in Seyfert 2 galaxies, their extended starburst activity is energetically more important than their compact nuclear starbursts. This is actually the case for the four Seyfert 2 galaxies studied in detail in the UV \\citep{gon98} and for the famous, well-studied Seyfert 2 galaxy NGC 1068 \\citep{lef01}. \\subsection{Do More Powerful AGNs Have More Powerful Compact Nuclear Starbursts?} We test the hypothesis that more powerful AGNs might be related to more powerful compact nuclear starbursts \\citep{gon98}. Figure~\\ref{fig2} compares the 12 $\\mu$m luminosity with the 3.3 $\\mu$m PAH emission luminosity. The total galactic 12 $\\mu$m luminosity is regarded as a measure of AGN power \\citep{gon01}. The 3.3 $\\mu$m PAH luminosities, measured with our slit spectroscopy, reflect the magnitudes of compact nuclear starbursts. The picture of \\citet{gon98} would be supported if we were to find a positive correlation between 12 $\\mu$m luminosity and 3.3 $\\mu$m PAH emission luminosity. In Fig.~\\ref{fig2}, this trend appears to be present, although the scatter is moderately large. We apply the generalized Kendall's rank correlation statistic \\citep{iso86} to the data points in Fig.~\\ref{fig2}, and estimate the probability that a correlation is not present to be 0.11, by using the software available at the web page: http://www.astro.psu.edu/statcodes/. Thus, provided that the 12 $\\mu$m luminosity is a good measure of AGN power, {\\it it is found that the luminosities of AGNs and compact nuclear starbursts in Seyfert 2 galaxies are correlated, and more powerful AGNs tend to contain more powerful compact nuclear starbursts.}" + }, + "0112/astro-ph0112536_arXiv.txt": { + "abstract": "P and CP violation in cosmology can be manifested as large-scale helical velocity flows in the ambient plasma and as primordial helical magnetic fields. We show that kinetic helicity at last scattering leads to temperature-polarization correlations ($C_l^{TB}$ and $C_l^{EB}$) in the cosmic microwave background radiation (CMBR) and calculate the magnitude of the effect. Helical primordial magnetic fields, expected from cosmic events such as electroweak baryogenesis, can lead to helical velocity flows and hence to non-vanishing correlations of the temperature and B-type polarization. However we show that the magnitude of the induced helical flow is unobservably small because the helical component of a magnetic field is almost force-free. We discuss an alternate scheme for extracting the helicity of a stochastically homogeneous and isotropic primordial magnetic field using observations of the CMBR. The scheme involves constructing Faraday rotation measure maps of the CMBR and thus determining the sum of the helical and non-helical components of the primordial magnetic field. The power spectrum of B-type polarization fluctuations, on the other hand, are sensitive only to the non-helical component of the primordial magnetic field. The primordial magnetic helicity can then be derived by combining these two sets of observations. ", + "introduction": "\\label{hvel} Kinetic helicity of a velocity field ${\\bf v}$ is characterized by a non-vanishing value of $\\langle {\\bf v}\\cdot {\\bm \\omega}\\rangle$, where ${\\bm \\omega}\\equiv \\nabla\\times {\\bf v}$. For the purpose of calculating the effect of helical flows on the cosmic microwave background radiation, it is useful to decompose the velocity field at last scattering into gradient components $v^{(0)}$ and rotational components $v^{(\\pm 1)}$ as follows: \\begin{eqnarray} \\label{eq:vdecomp} \\nonumber v_{j}\\left( {\\mathbf x} \\right) &=& \\int {d^{3}k \\over (2\\pi)^3}\\, e^{-i{\\mathbf k}\\cdot {\\mathbf x}} \\\\ & \\times & \\left( i\\hat{k}_{j} v^{(0)}+Q_{j}^{(1)}v^{(1)}+ Q_{j}^{(-1)}v^{(-1)}\\right) , \\end{eqnarray} where \\begin{equation} {\\mathbf Q}^{(\\pm 1)}\\left( {\\mathbf k}\\right) \\equiv -i\\frac{\\hat{{\\mathbf e}}_{1}\\pm i\\hat{{\\mathbf e}}_{2}}{\\sqrt{2}} \\end{equation} and the vectors $\\hat{{\\mathbf e}}_{1,2}$ are such functions of ${\\mathbf k}$ that they form an orthonormal basis together with $\\hat{{\\mathbf e}}_{3}\\equiv \\hat{{\\mathbf k}}$. The functions ${\\mathbf Q}^{(\\pm 1)}$ can be thought of as eigenvectors of helicity because $i\\hat{{\\mathbf k}}\\times {\\mathbf Q}^{(s)}=s{\\mathbf Q}^{(s)}$, where $s=\\pm 1$. Also, ${\\mathbf k}\\cdot {\\mathbf Q}^{(s)}=0$ and ${\\mathbf Q}^{(s)}\\left( {\\mathbf k}\\right) ={\\mathbf Q}^{(-s)}\\left( -{\\mathbf k}\\right) $. Under a parity transformation, ${\\mathbf v}({\\mathbf x})\\rightarrow -{\\mathbf v}({\\mathbf {-x}})$, ${\\mathbf k}\\rightarrow -{\\mathbf k}$ while ${\\mathbf Q}^{(\\pm 1)}\\rightarrow {\\mathbf Q}^{(\\mp 1)}$ and therefore $v^{(0)}\\left( {\\mathbf k}\\right) \\rightarrow -v^{(0)}\\left( -{\\mathbf k}\\right) $, $v^{(\\pm 1)}\\left( {\\mathbf k}\\right) \\rightarrow -v^{(\\mp 1)}\\left( -{\\mathbf k}\\right) $. In Appendix~\\ref{appendixA} we will show that the parity-odd CMB correlators $C_{l}^{TB}$ and $C_{l}^{EB}$ are linearly dependent on the expectation value of the parity-odd combination of $v^{(\\pm 1)}$: \\begin{equation} \\langle v^{(1)}\\left( {\\mathbf k}\\right) v^{(1)}\\left( -{\\mathbf k}\\right) -v^{(-1)}\\left( {\\mathbf k}\\right) v^{(-1)} \\left( -{\\mathbf k}\\right)\\rangle \\, . \\label{eq:vv} \\end{equation} The average helicity of the velocity field is proportional to the same quadratic combination of $v^{(\\pm 1)}$. Namely, \\begin{eqnarray} &&\\int d^{3}{\\bf x} ~ {\\mathbf v}( {\\mathbf x}+{\\mathbf y} ) \\cdot [\\nabla \\times {\\mathbf v} ( {\\mathbf x})] = \\int {d^{3}{\\mathbf k} \\over (2\\pi)^3} k e^{i{\\mathbf k}\\cdot {\\mathbf y}} \\nonumber \\\\ &&\\times \\left( v^{(1)}( {\\mathbf k}) v^{(1)}( -{\\mathbf k}) - v^{(-1)}( {\\mathbf k}) v^{(-1)}( -{\\mathbf k}) \\right), \\end{eqnarray} where $k\\equiv |{\\bf k}|$. Let us assume that the initial velocity field is random with a given power spectrum, such as \\begin{equation} \\left\\langle v_{i}( {\\mathbf k}) v_{j}( {\\mathbf k}') \\right\\rangle = v_{0}^{2} \\frac{k^n}{k_{*}^{n+3}} (2\\pi)^3 \\delta \\left( {\\mathbf k}+{\\mathbf k}'\\right) P_{ij}( \\hat{{\\mathbf k}}) , \\label{vikvjk} \\end{equation} where $v_0$ is the characteristic velocity of the flow and $k_*$ is the wave vector corresponding to a cutoff scale. In general, $P_{ij}$ does not have to be a symmetric tensor. The restrictions on it are: the reality condition $P^{*}_{ij}( \\hat{{\\mathbf k}}) = P_{ij}(-\\hat{{\\mathbf k}}) = P_{ji}( \\hat{{\\mathbf k}})$, the divergence-less condition (incompressibility) $k^{i}P_{ij}=0$, and the requirement of correct transformation under rotations, which forces $P_{ij}( \\hat{{\\mathbf k}})$ to be a tensorial function of $\\hat{k}^i$. The latter condition is needed to ensure that the correlator $P_{ij}(\\hat{{\\mathbf k}})$ describes a homogeneous and isotropic random vector field. It follows from these restrictions that \\begin{equation} \\label{eq:Pij} P_{ij}( \\hat{{\\mathbf k}}) = f(k) \\left[ \\delta _{ij}-\\hat{k}_{i}\\hat{k}_{j}\\right] + ig(k) \\varepsilon_{ijl}\\hat{k}^{l}, \\label{Pijfg} \\end{equation} where $f(k)$ and $g(k) $ are real functions of $k=|{\\bf k}|$. By using $$ \\left \\langle |{\\bm \\omega}({\\bf k}) \\pm k {\\bf v}({\\bf k})|^2 \\right \\rangle \\ge 0, $$ where ${\\bm \\omega}({\\bf k})$ is the Fourier transform of ${\\bm \\omega}({\\bf x})$, we can deduce that the functions $f$ and $g$ are not completely independent but must satisfy the inequality \\begin{equation} f({\\bf k}) \\ge |g({\\bf k})| \\, . \\label{maxkinhel} \\end{equation} [$f({\\bf k}) \\ge 0$ follows by taking the trace of Eq.~(\\ref{vikvjk}).] Eqs.~(\\ref{eq:vdecomp}) and (\\ref{vikvjk}) now lead to \\begin{eqnarray} \\left\\langle v^{(1)}\\left( {\\mathbf k}\\right) v^{(1)}\\left( {\\mathbf k}'\\right) -v^{(-1)}\\left( {\\mathbf k}\\right) v^{(-1)}\\left( {\\mathbf k}'\\right) \\right\\rangle \\nonumber \\\\ = v_{0}^{2} \\frac{k^n}{k_{*}^{n+3}} (2\\pi)^3 \\delta \\left( {\\mathbf k}+{\\mathbf k}'\\right) P^{ij}( \\hat{{\\mathbf k}}) Q_{[i}^{(1)}( \\hat{{\\mathbf k}}) Q_{j]}^{(-1)}( \\hat{{\\mathbf k}}) . \\end{eqnarray} The $f$-term in Eq.~(\\ref{Pijfg}) will not produce helicity because it is symmetric in $i$, $j$. Since \\begin{equation} i\\varepsilon _{ijl}\\hat{k}^{l}Q_{[i}^{(1)}( \\hat{{\\mathbf k}}) Q_{j]}^{(-1)}( \\hat{{\\mathbf k}}) =1, \\end{equation} we get \\begin{eqnarray} \\label{odd_{s}pectrum} \\left\\langle v^{(1)}( {\\mathbf k})v^{(1)}( {\\mathbf k}') - v^{(-1)}( {\\mathbf k}) v^{(-1)}( {\\mathbf k}') \\right\\rangle \\nonumber \\\\ = v_{0}^{2} \\frac{k^n}{k_{*}^{n+3}} (2\\pi)^3 \\delta \\left( {\\mathbf k}+{\\mathbf k}'\\right) g(k) . \\end{eqnarray} ", + "conclusions": "\\label{conclusions} We have analyzed certain P- and CP-violating signatures in the CMBR. If there is kinetic helicity at last scattering, it would imprint a signature in the cross-correlators $C_l^{TB}$ and $C_l^{EB}$. Kinetic helicity can be induced by helical magnetic fields but the effect is too small to be significant since the helical component of magnetic fields is force-free. Instead we have proposed another strategy for detecting the helicity of primordial magnetic fields using polarization and rotation measure maps of the CMBR. The helical magnetic fields produced during electroweak baryogenesis ($\\sim 10^{-13}$~G at last scattering) are several orders weaker than current upper bounds on the magnetic field strength ($\\sim 10^{-6}$~G) at last scattering \\cite{Vac01a}. Therefore the detection of electroweak fields does not seem feasible with forthcoming experiments. However, it is conceivable that stronger helical fields were produced due to some other mechanism and so it is still important to think of strategies for detecting primordial magnetic fields and helicity." + }, + "0112/astro-ph0112193_arXiv.txt": { + "abstract": "We present a measurement of the absolute surface brightness of the zodiacal light (3900--5100\\AA) toward a fixed extragalactic target at high ecliptic latitude based on moderate resolution ($\\sim$1.3\\AA\\ per pixel) spectrophotometry obtained with the du Pont 2.5m telescope at Las Campanas Observatory in Chile. This measurement and contemporaneous Hubble Space Telescope data from WFPC2 and FOS comprise a coordinated program to measure the mean flux of the diffuse extragalactic background light (EBL). The zodiacal light at optical wavelengths results from scattering by interplanetary dust, so that the zodiacal light flux toward any extragalactic target varies seasonally with the position of the Earth. This measurement of zodiacal light is therefore relevant to the specific observations (date and target field) under discussion. To obtain this result, we have developed a technique that uses the strength of the zodiacal Fraunhofer lines to identify the absolute flux of the zodiacal light in the multiple--component night sky spectrum. Statistical uncertainties in the result are 0.6\\% ($1\\sigma$). However, the dominant source of uncertainty is systematic errors, which we estimate to be 1.1\\% ($1\\sigma$). We discuss the contributions included in this estimate explicitly. The systematic errors in this result contribute 25\\% in quadrature to the final error in our coordinated EBL measurement, which is presented in the first paper of this series. ", + "introduction": "\\label{intro} This is the second in a series of three papers in which we present a measurement of the mean flux of the optical extragalactic background light and the cosmological implications of that result (see Bernstein, Freedman, \\& Madore 2002a \\& 2002c). The extragalactic background light (EBL) is the spatially averaged surface brightness of all extragalactic sources, resolved and unresolved. As such, the absolute flux of the EBL is a powerful and fundamental cosmological constant which can significantly constrain galaxy formation and evolution scenarios. Like all diffuse backgrounds, however, the optical EBL is very difficult to isolate from foreground sources, which are two orders of magnitude brighter. At high Galactic and ecliptic latitudes ($>30^\\circ$), the sky flux observed from the ground is dominated by terrestrial airglow and zodiacal light (ZL), each with a surface brightness of $\\sim 23$ AB mag arcsec$^{-2}$. The Hubble Space Telescope (HST), which orbits at an altitude of 600 km, avoids atmospheric emission, but the total sky flux is still dominated by ZL. An accurate measurement of the ZL is therefore crucial to a successful detection of the diffuse EBL. Our measurement of the EBL involves simultaneous HST and ground--based observations. From HST we measure the total flux of the night sky, including ZL. Using spectrophotometry over the range 3860--5150\\AA\\ (1.25\\AA\\ per pixel) taken with the Boller \\& Chivens Spectrograph on the duPont 2.5m Telescope at Las Campanas Observatory in Chile, we measure the absolute flux of the ZL contributing to the HST observations, which we can then subtract. In Bernstein, Freedman, \\& Madore (2002a, henceforth Paper I), we present the full details of the coordinated program to measure the EBL. In this paper, we present the ground--based measurement of the absolute flux of the ZL. As calibration of these data is crucial to the scientific goals, the data acquisition, reduction, and flux calibration are discussed here in detail. Background regarding the nature of the ZL is given in \\S\\ref{backg}. The observations, data reduction, and flux calibration are discussed in \\S\\ref{lco.obser}. In \\S\\ref{nightsky}, we briefly described the complications which arise due to atmospheric scattering, which redirects off--axis flux into and on--axis flux out of the line of sight. Detailed calculations of the atmospheric scattering relevant to precisely our observing situation (defined by the observatory location and positions of the Sun, Galaxy, and target relative to eachother and the horizon) are relegated to the Appendix, and summarized in \\S\\ref{nightsky}. In \\S\\ref{lco.analy}, we describe the technique used to measure the zodiacal light flux in reduced spectra. The results are summarized in \\S\\ref{lco.resul}. ", + "conclusions": "\\label{lco.discuss} \\begin{figure}[t] \\begin{center} \\includegraphics[width=3in,angle=00]{bfm2_fig13.ps} \\caption{\\footnotesize The same as Figure \\ref{fig:lco.results2}, but here only the four features with the smallest error bars have been used to produce the mean value of the ZL at 4600--4700\\AA\\ in each spectrum The horizontal line shows the mean ZL flux, which is $109.4(\\pm0.6)\\times10^{-9}$ \\escsa. Dashed lines show the one-sigma statistical error in the mean (0.6\\%). } \\label{fig:lco.results3} \\end{center} \\end{figure} \\begin{figure*}[t] \\begin{center} \\includegraphics[width=5in,angle=-90]{bfm2_fig14.ps} \\caption{\\footnotesize Airglow emission lines in the observed spectrum of the night sky after zodiacal light is subtracted. Broad emission features (marked above) are molecular rotation--vibration bands of O$_2$, N$_2$, H$_2$, OH, and NO$_2$. Collisional de-excitation of these molecules contributes a continuum as well. Emission lines (marked below) are also seen in this wavelength range from atomic transitions (photoionized O and Hg). } \\label{fig:lco.airglow} \\end{center} \\end{figure*} The {\\it rms} scatter in our ZL solution is less than 1\\%. This demonstrates that statistical errors are quite small, be they a result of instrumental effects or our analysis method. Systematic uncertainties due to instrumental effects are also quite small and are straightforward to quantify (see Table \\ref{tab:lco.errorbudget}). This is demonstrated by the fact that we obtain consistent results to within 0.3\\% from data taken on two different nights, independently reduced and calibrated. The uncertainties in the atmospheric scattering model described in the Appendix have been considered very carefully and we believe that the adopted uncertainties are conservative. However, the measurement presented here is obviously complex and might be effected by systematic errors which are more difficult to anticipate or quantify. One such systematic effect might include moonlight or sunlight scattered in the atmosphere. Based on the scattering analysis in the Appendix, it is clear that scattering into the line of sight from any source, even the sun or moon, is negligible when that source is more than 14 degrees below the horizon. All of the observations in this work took place when the sun was more than 18 degrees below the horizon. Although the observations took place several days after new moon, the moon was below the horizon during all of our program observations, and below 14 degrees for all but 1 exposure. In addition, the net effect of such solar-type scattering contributions, if present, would be to {\\it increase} our estimate of the zodiacal light, and consequentally to artificially {\\it decrease} the value of the inferred EBL in Paper I. We do not believe that such scattering is likely to have influenced our results. As discussed in \\S\\S \\ref{lco.analy} and \\ref{lco.resul}, the contamination of solar features by airglow, while introducing a systematic error, would not introduce a {\\it stable} systematic error: the flux of airglow features changes constantly through the night. The stability of our ZL solution in the eight spectral regions we have used demonstrates empirically that airglow is unlikely to have had a significant effect on our results. However, the possibility can't be ruled out and may introduce a systematic error which we cannot quantify and is not included in our estimate of the formal errors. It is also possible that some Doppler shifting occurs in the ZL spectral features relative to the solar spectrum due to motion of the dust in the zodiacal plane. For that reason, we allowed for a shift in the central wavelength when calculating the correlation but found no measurable offset. We note, also, that the results of this method would not be affected by the slight Doppler broadening which might affect the spectral features of the ZL, because Doppler broadening will not alter the total flux across a feature. The correlation is unaffected by the saw-tooth effect of subtracting features with mismatched widths at the level of the 0.3\\AA\\ Doppler broadening which is expected at the orientation of these observations (East \\& Reay 1984). Note also that while the resolution of the input spectra used for calculating the scattered ISL flux is lower than the resolution of our program observations (4\\AA\\ versus 2.6\\AA), this will not affect our analysis as long as regions with width $>>$4\\AA\\ are used in the analysis. The smallest of our spectral regions is 15\\AA. Figure \\ref{fig:lco.airglow} shows the airglow spectrum (the night sky spectrum after zodiacal light is subtracted) we obtain by this method. Emission lines from molecular rotation-vibration bands (O$_2$, N$_2$, H$_2$, OH, and NO$_2$) are labeled, as are some atomic transmission lines (O and Hg). Identification of emission features in this range of the spectrum is not complete (see Schmidtke \\etal 1985, Slanger \\& Huestis 1981, Jones \\etal 1985 and references therein). Finally, we note that our measurement of both the mean flux and ZL color are in very good agreement with typical values for the similar viewing geometries quoted in the literature (see the results of Levasseur--Regourd \\& Dumont 1980, pictured in Figure \\ref{fig:ZLoverSky}, and Leinert 1998)." + }, + "0112/astro-ph0112470_arXiv.txt": { + "abstract": " ", + "introduction": "The {\\it Midcourse Space Experiment} (\\msx) was a US Ballistic Missile Defense Organization (BMDO) satellite launched on 26 April 1996, from Vandenberg AFB into a 900 km altitude, semi-sun synchronous orbit. Mill et al. (1994) provide an overview of the \\msx\\ mission and objectives, which included observations of a variety of natural and man-made phenomena over a spectral range from the ultraviolet to the mid-infrared. A large number of astronomical observations were made during the mission, characterizing various components of celestial backgrounds. Price et al. (1997) describe the objectives of the 11 \\msx\\ astronomy experiments, designated ``CB\" (for Celestial Background), and also detail the data obtained by SPIRIT III, the infrared telescope aboard \\msx. The ultraviolet and visible sensor suite of four imagers and five hyperspectral imagers (collectively called UVISI) on \\msx\\ are described by Hefferman et al. (1996) and Paxton et al. (1996). In the first part of the mission (May, 1996 to February, 1997), the infrared instruments had the highest priority and were used to scan many interesting regions of the sky. Although the two UV imagers and five hyperspectral imagers (SPIMs) did take data during this phase, while the IR measurements were being made, it was only after the cryogen ran out (February 1997) that the full suite of UVISI sensors were used to begin a systematic survey of the sky as well as dedicated observations of specific regions. More than 400 observations (covering a substantial fraction of the sky) were obtained by the end of this phase of the mission in December 1997. These experiments were denoted as CB-10 (a systematic survey of the sky) and CB-11 (observations of specific celestial targets). The Orion nebula has long been known to be one of the brightest regions of diffuse UV emission in the sky since the first (and only) observations from a sounding rocket flight by Carruthers and Opal (1977). Despite considerable internal (instrumental) scattering, they detected a diffuse signal amounting to approximately 10\\% of the direct starlight, which they attributed to the scattering of starlight from interstellar dust in Orion. In the present paper, we describe the three {\\em MSX} observations, or ``data collection events'' (referred to as DCEs hereafter), in which the UVISI sensors were directed toward various regions in Orion. We detected intense UV emission in all three fields. In the following sections we describe our observations and results. ", + "conclusions": "We have observed intense diffuse radiation from three fields around M42 in Orion. This background is much brighter to the east and west of the nebula with intensities of more that 2 x 10$^4$ \\phunit\\ dropping to 8000 \\phunit\\ to the south of M42. In one of the DCEs (CB11-46 to the east of M42), we detected a bright patch which may be related to the nearby molecular clouds of Ogura \\& Sugitani (1998). In the other two regions, the emission was uniform. The UV flux was directly correlated with the IR in CB11-46 with a ratio of about 40 \\phunit\\ (MJy sr$^{-1}$)$^{-1}$. No such relation was observed in CB11-45 where, despite a factor of 4 variation in the \\iras\\ 100 \\micron\\ flux, the UV was essentially constant over the field. The only other measurement of the UV-IR ratio was the 128 \\phunit\\ (MJy sr$^{-1}$)$^{-1}$ obtained by Haikala et al (1995) for an isolated Galactic cirrus cloud. It is clear, and should be expected, that the UV-IR ratio is heavily dependent on the environment. In principle, the combination of IR and UV observations in a single direction will strongly constrain the optical properties of the dust grains. Stellar radiation penetrates into the interior of the clouds and some part is scattered in the UV. That part not scattered will be absorbed and then reemitted as thermal radiation in the IR. In practice, a detailed study is required to realistically model the penetration and subsequent reradiation of the input radiation field by the optically thick (in the UV) dust. We are beginning such a study for simple geometries where we have Voyager and IRAS data and hope, eventually, to extend it to more complex regions such as Orion. This work was supported at the Johns Hopkins University by USAFG F19628-93-K-0004" + }, + "0112/nucl-th0112002_arXiv.txt": { + "abstract": "s#1#2#3{{ \\centering{\\begin{minipage}{4.5in}\\footnotesize\\baselineskip=10pt \\parindent=0pt #1\\par \\parindent=15pt #2\\par \\parindent=15pt #3 \\end{minipage}}\\par}} \\def\\keywords#1{{ \\centering{\\begin{minipage}{4.5in}\\footnotesize\\baselineskip=10pt {\\footnotesize\\it Keywords}\\/: #1 \\end{minipage}}\\par}} \\newcommand{\\bibit}{\\nineit} \\newcommand{\\bibbf}{\\ninebf} \\renewenvironment{thebibliography}[1] {\\frenchspacing \\ninerm\\baselineskip=11pt \\begin{list}{\\arabic{enumi}.} {\\usecounter{enumi}\\setlength{\\parsep}{0pt} \\setlength{\\leftmargin 12.7pt}{\\rightmargin 0pt} % \\setlength{\\itemsep}{0pt} \\settowidth {\\labelwidth}{#1.}\\sloppy}}{\\end{list}} \\newcounter{itemlistc} \\newcounter{romanlistc} \\newcounter{alphlistc} \\newcounter{arabiclistc} \\newenvironment{itemlist} {\\setcounter{itemlistc}{0} \\begin{list}{$\\bullet$} {\\usecounter{itemlistc} \\setlength{\\parsep}{0pt} \\setlength{\\itemsep}{0pt}}}{\\end{list}} \\newenvironment{romanlist} {\\setcounter{romanlistc}{0} \\begin{list}{$($\\roman{romanlistc}$)$} {\\usecounter{romanlistc} \\setlength{\\parsep}{0pt} \\setlength{\\itemsep}{0pt}}}{\\end{list}} \\newenvironment{alphlist} {\\setcounter{alphlistc}{0} \\begin{list}{$($\\alph{alphlistc}$)$} {\\usecounter{alphlistc} \\setlength{\\parsep}{0pt} \\setlength{\\itemsep}{0pt}}}{\\end{list}} \\newenvironment{arabiclist} {\\setcounter{arabiclistc}{0} \\begin{list}{\\arabic{arabiclistc}} {\\usecounter{arabiclistc} \\setlength{\\parsep}{0pt} \\setlength{\\itemsep}{0pt}}}{\\end{list}} \\newcommand{\\fcaption}[1]{ \\refstepcounter{figure} \\setbox\\@tempboxa = \\hbox{\\footnotesize Fig.~\\thefigure. #1} \\ifdim \\wd\\@tempboxa > 5in {\\begin{center} \\parbox{5in}{\\footnotesize\\smalllineskip Fig.~\\thefigure. #1} \\end{center}} \\else {\\begin{center} {\\footnotesize Fig.~\\thefigure. #1} \\end{center}} \\fi} \\newcommand{\\tcaption}[1]{ \\refstepcounter{table} \\setbox\\@tempboxa = \\hbox{\\footnotesize Table~\\thetable. #1} \\ifdim \\wd\\@tempboxa > 5in {\\begin{center} \\parbox{5in}{\\footnotesize\\smalllineskip Table~\\thetable. #1} \\end{center}} \\else {\\begin{center} {\\footnotesize Table~\\thetable. #1} \\end{center}} \\fi} \\def\\@citex[#1]#2{\\if@filesw\\immediate\\write\\@auxout {\\string\\citation{#2}}\\fi \\def\\@citea{}\\@cite{\\@for\\@citeb:=#2\\do {\\@citea\\def\\@citea{,}\\@ifundefined {b@\\@citeb}{{\\bf ?}\\@warning {Citation `\\@citeb' on page \\thepage \\space undefined}} {\\csname b@\\@citeb\\endcsname}}}{#1}} \\newif\\if@cghi \\def\\cite{\\@cghitrue\\@ifnextchar [{\\@tempswatrue \\@citex}{\\@tempswafalse\\@citex[]}} \\def\\citelow{\\@cghifalse\\@ifnextchar [{\\@tempswatrue \\@citex}{\\@tempswafalse\\@citex[]}} \\def\\@cite#1#2{{$\\null^{#1}$\\if@tempswa\\typeout {IJCGA warning: optional citation argument ignored: `#2'} \\fi}} \\newcommand{\\citeup}{\\cite} \\def\\pmb#1{\\setbox0=\\hbox{#1} \\kern-.025em\\copy0\\kern-\\wd0 \\kern.05em\\copy0\\kern-\\wd0 \\kern-.025em\\raise.0433em\\box0} \\def\\mbi#1{{\\pmb{\\mbox{\\scriptsize ${#1}$}}}} \\def\\mbr#1{{\\pmb{\\mbox{\\scriptsize{#1}}}}} \\def\\fnm#1{$^{\\mbox{\\scriptsize #1}}$} \\def\\fnt#1#2{\\footnotetext{\\kern-.3em {$^{\\mbox{\\scriptsize #1}}$}{#2}}} \\def\\fpage#1{\\begingroup \\voffset=.3in \\thispagestyle{empty}\\begin{table}[b]\\centerline{\\footnotesize #1} \\end{table}\\endgroup} \\def\\thefootnote{\\fnsymbol{footnote}} \\def\\@makefnmark{\\hbox to 0pt{$^{\\@thefnmark}$\\hss}}\t% \\def\\ps@myheadings{% \\let\\@oddfoot\\@empty\\let\\@evenfoot\\@empty \\def\\@evenhead{\\slshape\\leftmark\\hfil}% \\def\\@oddhead{\\hfil{\\slshape\\rightmark}}% \\let\\@mkboth\\@gobbletwo \\let ", + "introduction": " ", + "conclusions": "" + }, + "0112/hep-th0112115_arXiv.txt": { + "abstract": "Four dimensional gravity in the low energy limit of a higher dimensional theory has been expected to be a (generalized) Brans-Dicke theory. A subtle point in brane world scenarios is that the system of four dimensional effective gravitational equations is not closed due to bulk gravitational waves and bulk scalars. Nonetheless, weak gravity on the brane can be analyzed completely. We revisit the theory of weak brane gravity using gauge-invariant gravitational and scalar perturbations around a background warped geometry with a bulk scalar between two flat branes. We obtain a simple condition for the radion stabilization in terms of the scalar field potentials. We show that for general potentials of the scalar field which provides radion stabilization and a general conformal transformation to a frame in which matter on the branes are minimally coupled to the metric, $4$-dimensional Einstein gravity, not BD gravity, is restored at low energies on either brane. In contrast, in RS brane world scenario without a bulk scalar, low energy gravity is BD one. We conjecture that in general brane world scenarios with more than one scalar field, one will again encounter the situation that low energy gravity is not described by the Einstein theory. Equipped with the weak gravity results, we discuss the properties of 4d brane gravitational equations, in particular, the value and sign of 4d Newton's gravitational coupling. ", + "introduction": "In the recent development of string/M theory~\\cite{Polchinski}, branes have been playing many important roles. The idea that our universe is a brane in a higher dimensional spacetime has been attracting a great deal of interest. Although the idea of the brane world had arisen at a phenomenological level already in 1983~\\cite{earlier-works}, it is perhaps the discovery of the duality between M-theory and $E_8\\times E_8$ heterotic superstring theory by Horava and Witten~\\cite{Horava} that made it more attractive. It actually gives the brane world idea a theoretical background: by compactifying six dimensions in the $11$-dimensional theory, our $4$-dimensional universe may be realized as a hypersurface in $5$-dimension at one of the fixed points of a $S_1/Z_2$. The $5$-dimensional effective theory after compactification of the six dimensions by a Calabi-Yau manifold can be obtained, e.g. \\cite{Lukas}. Randall and Sundrum proposed two similar but distinct phenomenological brane world scenarios~\\cite{RS1,RS2}. In both scenarios the $5$-dimensional spacetime is compactified on $S_1/Z_2$ and all matter fields are assumed to be confined on branes at fixed points of the $S_1/Z_2$ so that the bulk, or the spacetime region between two fixed points, is described by pure Einstein gravity with a negative cosmological constant. The first scenario~\\cite{RS1} deals with a three-brane with negative tension as our universe. It was shown that the hierarchy problem may be solved in this scenario by a large redshift (warp) factor, between our brane and another, hidden brane. In the second scenario~\\cite{RS2} a three-brane with positive tension is considered as our universe and four-dimensional Newtonian gravity can be realized on the brane in the limit of infinite orbifold radius. The distance between two branes (radion) in the phenomenological RS scenario, as well as, more generally, moduli fields in superstring theory, should be stabilized. For this purpose a bulk scalar field $\\Psi$ with certain potentials in the bulk and on the branes can be introduced~\\cite{GW}. The choice of the bulk $U(\\Psi)$ and brane potentials $V_{\\pm}(\\Psi)$ must be consistent with the 5d warp geometry \\cite{Dewolfe,GKL}. In the context of the supergravity realization of the RS scenario one has to consider bulk scalars of the theory with an effective scalar field potential~\\cite{super1,super2}. Thus, there is a broad class of brane world scenarios with a 5d bulk between two end-of-the-world branes and a bulk scalar field with bulk and brane potentials. This is a basis for a brane cosmology construction in which our 4d expanding universe is located on the visible brane. It is expected on the general ground that dimensional reduction of a multi-dimensional theory in the 4d low energy limit yields a (generalized) Brans-Dicke (BD) theory. Often 4d effective theory is derived by the integration of the action with respect to the extra dimension, $\\int dw$, as it customary in Kaluza-Klein compactification. In the brane world scenario one can derive effective 4d dimensional Einstein equations on the brane, starting with the full 5d bulk equations. Although at the qualitative level the expectation that you will obtain an effective BD theory is justified, many quantitative details are lost. The integration $\\int dw$ must be treated with special caution in cosmological models with brane collisions~\\cite{KKLT}. Widespread expectation is that BD scalar is a manifestation of a massless bulk scalar degree of freedom. If moduli are stabilized and no massless scalars left in the bulk, an expectation would be that BD theory is reduced to the pure Einstein gravity on the brane. Let us inspect a resent progress in respect to these ideas. ", + "conclusions": "\\label{sec:summary} We have investigated gravity in brane worlds. In particular, we considered a general model on a $Z_2$-orbifold with a bulk scalar field and investigated gauge-invariant perturbations around a background with two parallel Minkowski branes. By using the doubly covariant formalism of perturbation of Israel's junction conditions~\\cite{Israel} developed in ref.~\\cite{Mukohyama2000b}, we have obtained equations governing the low energy dynamics of doubly gauge invariant perturbations on the branes. These doubly gauge invariant perturbations include scalar, vector, and tensor perturbations. After that, we have shown that $4$-dimensional Einstein gravity is restored on both branes at low energies, provided that combinations of the bulk and the brane potentials $\\tilde{B}_{\\pm}$ given by (\\ref{eqn:BandtildeB}) are not zero at the branes and that there are no tachyonic modes for scalar perturbations with the boundary condition (\\ref{eqn:tau=0-bc}). It is easy to see that the condition $\\tilde{B}_{\\pm}\\ne 0$ is equivalent to the absence of a scalar zero mode with vanishing matter on the branes. Actually, the would-be zero mode with vanishing matter is prohibited by the boundary condition (\\ref{eqn:C1-C2-bc}) with $\\bar{\\tau}_{\\pm(LL)}=0$ if and only if $\\tilde{B}_{\\pm}=0$. The effective $4$-dimensional Newton's constants $G_{N\\pm}$ on $\\Sigma_{\\pm}$, respectively, are given by (\\ref{eqn:Newton-const}) and the naive expectation (\\ref{eqn:warp-factor}) holds. Thus we have confirmed and extended the result of \\cite{Tanaka-Montes} that pure Einstein gravity is restored for the braneworld scenario with a single bulk scalar at low energies. The condition $\\tilde{B}_{\\pm}\\ne 0$ is equivalent to the condition that the inter-brane distance (radion) is stabilized (to be precise, fixed) and that the background bulk scalar field has non-zero derivatives near the branes. To see this it is enough to differentiate the background scalar-field matching condition (the second of (\\ref{eqn:background-junction})) with respect to $w_{\\pm}$~\\cite{Mukohyama2001c}. In fact, the radion stabilization can be stated as \\begin{eqnarray} 0 \\neq e^{A_{\\pm}}\\frac{\\partial}{\\partial w_{\\pm}} \\left[\\left.e^A\\dot{\\Psi}^{(0)}\\right|_{w=w_{\\pm}} \\pm\\frac{1}{2}V'_{\\pm}(\\Psi^{(0)}_{\\pm})\\right] = {U'_{\\pm}}^{(0)} +\\frac{\\kappa_5^2}{3}V_{\\pm}^{(0)}{V'_{\\pm}}^{(0)} -\\frac{1}{4}{V'_{\\pm}}^{(0)}{V''_{\\pm}}^{(0)} \\label{eqn:stabilization} \\end{eqnarray} since the following equation is always satisfied because of the background equation. \\begin{eqnarray} \\frac{\\partial}{\\partial w_{\\pm}} \\left[ \\left.e^A\\dot{A}\\right|_{w=w_{\\pm}} \\pm\\frac{1}{6}\\kappa_5^2V_{\\pm}(\\Psi^{(0)}_{\\pm})\\right]=0. \\end{eqnarray} Essentially, the condition (\\ref{eqn:stabilization}) says that the position of the brane cannot be changed without changing the bulk solution. Hence, actually the condition $\\tilde{B}_{\\pm}\\ne 0$ is equivalent to the condition that $V^{'(0)}_{\\pm}\\ne 0$ and that the inter-brane distance is fixed. Since $V^{'(0)}_{\\pm}$ is related to $\\dot{\\Psi}^{(0)}$ by (\\ref{eqn:background-junction}) on branes, finally the condition $\\tilde{B}_{\\pm}\\ne 0$ is equivalent to the condition that the background bulk scalar field is changing near the branes and that the radion is stabilized. One might expect that it is sufficient but not necessary to fix positions of both branes in order to fix the distance between two branes. In fact, it is necessary to fix positions of both branes as in (\\ref{eqn:stabilization}). To see this, let us assume that the position of one brane can be changed without changing the bulk solution at all. (The position of the other brane can be either fixed or unfixed.) In this case it is impossible for the position of the unfixed brane to affect the position of the other brane since the only possible medium of the communication between the two branes would be the bulk but the bulk solution is independent of the position of the unfixed brane by assumption. Thus, the distance between the branes is not fixed in this case. In other words, the matching condition for the background with the $4$-dimensional Poincare symmetry is local. Although to derive Newton's gravitational couplings, we have used only matter source at the brane $\\Sigma_+$ (visible brane), it is instructive to write down the most generic weak gravity equations taking into account the contribution from the hidden brane, which enters the final results in a symmetric manner, see the last equations in subsections \\ref{subsec:scalar} and \\ref{subsec:tensor}. Suppose energy-momentum tensor at the visible brane is $T^+_{\\mu\\nu}$ and that on the hidden brane $\\Sigma_-$ is $T^-_{\\mu\\nu}$. Let us again, as in Sec. II, use linearized metric perturbations $h_{\\mu\\nu}$ at the visible brane. Similar to (\\ref{eqn:BD-linear-decomposed}), in most general case we have \\begin{eqnarray} \\Box \\bar{h}^{TT}_{\\mu\\nu} & = & -16 \\pi G^+_N T^{+TT}_{\\mu\\nu} -16 \\pi G^-_N T^{-TT}_{\\mu\\nu} \\ , \\nonumber\\\\ \\Box h & = & {16 \\pi G^+_N}(1-2B^+)T^+ + {16 \\pi G^-_N}(1-2B^-)T^- \\ .\\label{eqn:BD-linear-decomposed1} \\end{eqnarray} Therefore, in principle, we need four gravitational couplings, two for tensor modes and two for scalar modes, to describe the linearized gravity at the visible brane. We also learned that each of them are constructed from the fundamental 5d gravitational constant, $\\kappa_5$, and combinations of the warp factor and ``effective'' extra dimensional distance $l \\sim {\\int^{w_+}_{w_-}dw'e^{-3A(w')}}$. It turns out that the BD parameter, which characterizes the gravitational coupling for the scalar mode, depends on the bulk scalar sector. In terms of coefficient $B$ of Sec. II, $B^{\\pm}$ is none-zero for the case without bulk scalar, $B^{\\pm}$=0 for the case of a single scalar field. We conjecture, based on our derivation with a single scalar, that in general case of more than one bulk scalar, $B^{\\pm}$ again will be non-zero. It will be interesting to calculate it explicitly for next simplest case of two bulk scalars. Equation (\\ref{eqn:BD-linear-decomposed1}) may have interesting insights from point of view of superstring (Horava-Witten) phenomenology. Indeed, it says that there are, in general, four different Planckian masses! It also teaches us that derivation of 4d effective theory by means of $\\int dw$ integration will wash out the subtleties of the actual gravitational dynamics. Finally, from point of view of gravity and its application to astronomy, it gives us rather unusual theory with four gravitational couplings." + }, + "0112/astro-ph0112464_arXiv.txt": { + "abstract": "{ The recent {\\em Chandra} observation of the radio source at the center of our Galaxy, Sgr A$^*$, puts new constraints on its theoretical models. The spectrum is very soft, and the source is rapidly variable. We consider different models to explain the observations. We find that the features of the x-ray spectrum can be marginally explained with an advection-dominated accretion flow (ADAF) model while it does not well fit the radio spectrum. An ADAF with strong winds (ADIOS) model is not favored if we assume that the wind does not radiate. Alternatively, we propose a coupled jet plus accretion disk model to explain the observations for Sgr A$^*$. The accretion flow is described as an ADAF fed by Bondi-Hoyle accretion of hot plasma in the Galactic Center region. A small fraction of the accretion flow is ejected near the black hole, forming a jet after passing through a shock. As a result, the electron temperature increases to $\\sim 2 \\times 10^{11}{\\rm K}$, which is about 10 times higher than the highest temperature attained in the ADAF. The model is self-consistent since the main jet parameters are determined by the underlying accretion disk at the inner edge. The emergent spectrum of Sgr A$^*$ is the sum of the emission from jet and underlying ADAF. The very strong Comptonization of synchrotron emission from the jet dominates the bremsstrahlung from the ADAF, therefore, a very short variability timescale is expected and the predicted X-ray slope and the radio spectrum is in very good agreement with the observations. ", + "introduction": "The energetic radio source Sgr A$^*$ located at the center of our Galaxy is now widely believed to be the signature of a massive black hole with mass $M=2.6 \\times 10^6 \\msun$ (Melia \\& Falcke 2001; Haller et al. 1996; Eckart \\& Genzel 1996; Ghez et al. 1998; Reid et al. 1999; Backer \\& Sramek 1999). Its radio spectrum seems to consist of two components, with a break around $\\sim 50$ GHz. The spectral dependence is $F_{\\nu} \\propto \\nu^{0.2}$ for $\\nu < 50 $ GHz, while above this break there is a submm bump which is described by $F_{\\nu} \\propto \\nu^{0.8}$ up to $\\sim 10^3$ GHz followed by a steep cut-off towards the infrared (IR) (Zylka et al. 1992; Serabyn et al. 1997; Falcke et al. 1998). The upper limits from IR (Menten et al. 1997) and ROSAT X-ray observations (Predehl \\& Tr\\\"{u}mper 1994) indicate that this source is quite dim. On the theoretical side, a number of models have been proposed in the past years for Sgr A$^*$. Most models are based on accretion onto the central massive black hole. Possible sources of accretion material include the stellar winds emitted by the nearby massive stars and the hot interstellar medium. Since in either case the angular momentum of the accretion flow should be small, Melia (1992; 1994) proposed a spherical accretion model. In this model the accretion flow is assumed to free-fall until a Keplerian disk is formed within a small ``circularization'' radius. The main contributors to the radio and X-ray spectra are synchrotron radiation and bremsstrahlung, respectively, from the roughly free-fall flow beyond the small disk. However, spherical accretion is likely to be an over-simplification, since the accretion flow still possesses some angular momentum. An advection-dominated accretion flow (ADAF) model therefore is more dynamically exact in this sense (Narayan et al. 1995; Manmoto et al. 1997; Narayan et al. 1998). The most attractive feature of the ADAF model is its ability to explain the unusual low-luminosity of Sgr A$^*$ given the relatively abundant accretion material. This is because most of the viscously dissipated energy is stored in the flow and advected beyond the event horizon rather than radiated away (Ichimaru 1977; Rees et al. 1982; Narayan \\& Yi 1994, 1995; Abramowicz et al. 1995; Chen et al. 1995; Narayan et al. 1997; Chen et al. 1997). In the application to Sgr A$^*$, the radio spectrum is produced by the synchrotron process in the innermost region of the disk while the X-rays are due to bremsstrahlung radiation of the thermal electrons in a large range of radii $\\sim 10^3-10^4 R_{\\rm s}$, where $R_{\\rm s} =2 GM/c^2$ is the Schwarzschild radius. However, the ADAF under-predicts the low-frequency radio emission of Sgr A$^*$ by over an order of magnitude and additional assumptions must be imposed in order to match the spectrum (Mahadevan 1998, \\\"{O}zel et al. 2000). Following the initial paper by Reynolds \\& McKee (1980) (see also Blandford \\& K\\\"onigl 1979), Falcke et al. (1993) proposed that it is the jet stemming from the disk rather than the disk itself which is responsible for the radio spectrum of Sgr A$^*$. In this model, the submm bump is produced by the acceleration zone of the jet, called nozzle, while the low-frequency radio spectrum comes from the part of the jet beyond the nozzle (Falcke 1996b; Falcke \\& Biermann 1999). The nozzle is of order 10$R_{\\rm s}$ and forms from the disk at a radius of $\\sim 2 R_{\\rm s}$. This model gives an excellent fit to the radio spectrum of Sgr A$^*$, including the low-frequency spectrum below the break and the submm bump, but the expected X-ray emission was not calculated explicitly. The latest observational constraints for Sgr A$^*$ come from the high spatial resolution ($\\approx 1^{\\arcsec}$) {\\em Chandra X-ray Observatory} (Baganoff et al. 2001a, 2001b). Baganoff et al. observed Sgr A$^*$ twice and they found that Sgr A$^*$ comes in two states: quiescent and flares. In the present paper we concentrate on the quiescent state, whereas the flare state is considered in Markoff et al. (2001b). The main observational results for the quiescent state are summarized as follows\\footnote{The luminosity and especially the photon index are taken from Baganoff et al. 2001b, which are slightly different from those in Baganoff et al. 2001a where the spectral models used did not account for dust scattering; see Baganoff et al. 2001b for details.}. \\begin{itemize} \\item The absorption-corrected 2-10 keV luminosity is (2.2$^{+0.4}_{-0.2}$)$\\times10^{33} {\\rm erg ~s}^{-1}$. \\item The spectrum is well fitted by an absorbed power-law model with photon index $\\Gamma=2.2^{+0.5}_{-0.7}$. \\item The inner region of the source is rapidly variable on short timescale of $\\simeq 1 {\\rm hr}$. A rapid drop of flux on a timescale of 10 minutes is detected in the flare state. On the other hand, the comparison between the two observations with an interval of about one year indicates that the steady X-ray flux remains almost constant. \\item Some fraction of the X-ray flux may come from a partly extended region with diameter $\\approx 1^{\\arcsec}$. \\item There is tentative evidence for a Fe K$\\alpha$ line at 6.7 keV. \\end{itemize} These results provide new and strict constraints to the theoretical models for Sgr A$^*$. In both the ADAF and spherical accretion models mentioned above, the X-ray radiation is produced by bremsstrahlung originating from $10^3$--$10^4R_{\\rm s}$. Hence the spectrum is very hard with photon index $\\Gamma \\sim 1.4$ and the predicted variability timescale is thousands of hours, much longer than the observed $\\sim 1$ hour. Therefore it is necessary to reexamine the theoretical models for Sgr A$^*$. Melia et al. (2001) proposed that the electrons in the small Keplerian disk can attain a very high temperature through some magnetic processes, and the resulting synchrotron and self-Compton emission are responsible for the radio and X-ray spectrum. However, the formation of the small disk may not be a necessary result of such low angular momentum accretion. An accretion flow with very low angular momentum can still be described by an ADAF, although such accretion may belong to the Bondi-like type rather than disk-like type, as shown by Yuan (1999) (see also Abramowicz \\& Zurek 1981; Abramowicz 1998). Thus the dynamical scenario of this model needs to be studied carefully. For the jet model, Falcke \\& Markoff (2000) take into account the contribution from synchrotron self-Compton emission (SSC) in the nozzle and find that the parameters required to interpret the submm bump give a very good fit to the {\\em Chandra} spectrum without changing the basic parameters of the jet model. But the remaining important problem in the model is why the parameters of the jet possess the required values, particularly in reference to the inferred underlying accretion disk. Previous ideas of a standard optically thick accretion disk in Sgr A* (e.g., Falcke \\& Heinrich 1994) do not seem to work because the predicted IR flux from a standard thin disk with a reasonable accretion rate would be several orders of magnitude higher than the observed IR upper limit (Falcke \\& Melia 1997). Therefore, it is crucial to consider the jet and accretion flow as a coupled system in Sgr A$^*$, and to consider what are their respective roles if both are truly present in Sgr A$^*$. Yuan (2000) presented the first effort, by considering a combination of jet and ADAF models. However, the complete {\\em Chandra} data was not available at that time and the detailed coupling mechanism was lacking in Yuan (2000) so it is necessary to revisit the model again. The development of the theory provides a new chance to model Sgr A$^*$. Since the Bernoulli parameter of the ADAF is positive, which means the gas can escape to infinity with positive energy, Blandford \\& Begelman (1999) propose an advection-dominated inflow-outflow solution (ADIOS) in which most of the gas is lost through winds rather than accreted past the horizon of the black hole. The concept of strong winds from accretion flow was also proposed and studied by Xu \\& Chen (1997) and Das \\& Chakrabarti (1999). The former described pressure-driven winds from centrifugally supported boundary layers and shocks in the inner regions of disks, and the latter proposed an advection-dominated flow where the central black hole redirects the inward flow at low latitudes into an outflow at high latitudes. We are not explicitly making use of the latter two models. The most appealing point of the ADIOS model as applied to Sgr A$^*$ is that the predicted X-ray spectrum is possibly much softer than that of the ADAF (Quataert \\& Narayan 1999), and therefore could possibly give a better fit to the {\\em Chandra} data. This is because the density profile of the accretion flow becomes flatter due to the wind, while X-ray emission at higher frequencies is produced in the inner region of the accretion flow. If we assume that the mass accretion rate in the ADIOS is described by a power-law of radius, $\\dot{M}\\propto R^{p}$, the predicted photon index in {\\em Chandra} band is approximately $\\Gamma \\approx 3/2+2p$. Thus it is necessary to investigate this model for the possibility of interpreting the {\\em Chandra} results. In this paper we explore several of the above-mentioned models for Sgr A$^*$. By probing a larger parameter space than before, we find that ADAFs can give a marginal interpretation to the new {\\em Chandra} results, although the fit is not very good in some points (Sect. 2), while the ADIOS model can't (Sect. 3). In Sect. 4 we propose that the combination of an ADAF and a jet could provide an excellent fit to the observations to Sgr A$^*$, and present our model results. The last section is a summary and discussion. ", + "conclusions": "Recent {\\em Chandra} X-ray observations put new constraints on the theoretical models of Sgr A$^*$. The spectrum is very soft, the flux is rapidly variable and the source is extended. In this paper we consider three different models to explain the observational results of Sgr A$^*$. We find that an ADAF model can give a marginally satisfactory interpretation to the {\\em Chandra} spectrum and the rapid X-ray variability. But our best fit is still not good for the radio spectrum in the sense that it over-predicts the high-frequency radio by a factor of 2-3 and significantly under-predicts the low-frequency radio. We then consider the possibility of strong winds from ADAFs, i.e., an ADIOS model. If the winds are non-radiative and viscous dissipation in the accretion flow mainly heats ions, as generally assumed in the literature, this model can fit the spectrum ranging from submm bump to X-ray quite well. However, it is hard to explain the rapid X-ray variability since in this model bremsstrahlung is the sole contributor at X-ray band. If we assume that most of the viscous dissipation preferentially heats electrons, a rapidly variable X-ray spectrum is expected since in this case the X-ray emission is dominated by SSC. But in this case the model over-predicts the radio flux above $\\sim 100$ GHz by a factor of 4-6, and the predicted X-ray spectrum is much steeper than the best fit of the {\\em Chandra} observations. An excellent fit to all the data including low-frequency radio can be obtained with a coupled jet-disk model. In this model, the accretion disk is described by an ADAF. In the innermost region of the ADAF, $\\sim 2R_{\\rm s}$, some fraction $q_{\\rm m}$ ($\\sim 0.5\\%$ if any cold jet component is neglected. See our discussion in Sect. 4 for the possibility of a cold jet component) of the accretion flow is ejected out of the ADAF and transferred into the jet. In this process, a shock occurs because the accretion flow is radially supersonic before the shock. After the shock the temperature of electrons in the nozzle (the base of the jet) reaches about $2 \\times 10^{11}{\\rm K}$. In this case, the synchrotron emission in the nozzle largely dominates the submm bump, and its Comptonization dominates the quiescent X-ray spectrum in Sgr A$^*$. The X-ray spectrum is soft and the variability timescale is short. Out of the nozzle, the jet gas expands freely outward under the force of the gas pressure gradient of gas pressure. Furthermore its self-absorbed synchrotron radiation gives an good fit to the low-frequency radio spectrum of Sgr A$^*$ which is hard to explain in ADAF models. The model is completely self-consistent. The jet in our model produces a slightly inverted radio spectrum, as can be understood from the canonical model of Blandford \\& K\\\"onigl (1979), with modifications as in Falcke (1996a). In the absence of a shock acceleration region in the highly-supersonic outer region of the jet, the particles retain the highly-peaked relativistic Maxwellian energy distribution which is attained by shock {\\it heating} occurring when the radial supersonic accretion flow is transferred into the vertical direction. On the other hand, the electrons in AGN jets typically seem to have a power-law high-energy tail after shock acceleration in jets, since the Mach number in jets is very high (Drury 1983). In that case, a corresponding optically thin power-law spectrum at IR/optical frequencies is generally expected, as is seen in many AGN and perhaps even X-ray binary jets (e.g., Markoff et al. 2001a). In the case of Sgr A$^*$, the absence of an optically-thin power-law indicates that, for some unknown reason, such high Mach number shocks do not occur. If they would occur under certain conditions, we should still see an inverted radio spectrum, but we would also expect some kind of hard power-law emission at higher frequencies (mid-IR to X-rays). In addition to the observations we mention in the present paper, there are also constraints to the model through the frequency-size relationship obtained from VLBI observations (Rogers et al. 1994; Krichbaum et al. 1998; Lo et al. 1998). The jet-disk model can fit this well as shown in Falcke \\& Markoff (2000). We therefore conclude that it is possible to present a consistent picture of the emission processes associated with the central black hole in our Galaxy by combining the three basic astrophysical ingredients that have been discussed in recent years: Bondi-Hoyle accretion from the immediate environment, optically thin accretion through an ADAF, and energy extraction and visible emission by a plasma jet. Our jet-ADAF model predicts a closely correlated variability among sub-millimeter, IR, and X-ray. More broad-band observations and monitoring at various wavebands (radio, IR, X-rays) will help to judge whether it will be possible to establish a standard model invoking those elements for Sgr A* in the near future. For example, more precise determination of the IR flux will help to further discriminate between the jet-ADAF model and the pure ADAF model since the former predicts higher IR flux than the latter. This will also be crucial for understanding the activity in low-power black holes in general." + }, + "0112/astro-ph0112187_arXiv.txt": { + "abstract": "{ We present measurements of the emission from the centers of fifteen spiral galaxies in the $\\pci$ fine-structure transition at 492 GHz. Observed galaxy centers range from quiescent to starburst to active. The intensities of neutral carbon, the $J$=2--1 transition of $\\13co$ and the $J$=4--3 transition of $\\co$ are compared in matched beams. Most galaxy centers {\\it emit more strongly in [CI] than in $\\13co$}, completely unlike the situation pertaining to Galactic molecular cloud regions. [CI] intensities are lower than, but nevertheless comparable to $J$=4--3 $\\co$ intensities, again rather different from Galactic sources. The ratio of [CI] to $\\13co$ increases with the central [CI] luminosity of a galaxy; it is lowest for quiescent and mild starburst centers, and highest for strong starburst centers and active nuclei. Comparison with radiative transfer model calculations shows that most observed galaxy centers have {\\it neutral carbon abundances close to, or exceeding, carbon monoxide abundances}, rather independent from the assumed model gas parameters. The same models suggest that the emission from neutral carbon and carbon monoxide, if assumed to originate in the same volumes, {\\it arises from a warm and dense gas} rather than a hot and tenuous, or a cold and very dense gas. The observed [CI] intensities together with literature [CII] line and far-infrared continuum data likewise suggest that a significant fraction of the emission originates in medium-density gas ($n = 10^{3}-10^{4}\\cc$), subjected to radiation fields of various strengths. ", + "introduction": "\\begin{figure*}[t] \\unitlength1cm \\begin{minipage}[]{16.4cm} \\resizebox{17cm}{!}{\\rotatebox{270}{\\includegraphics*{cifig_1a.ps}}} \\end{minipage} \\caption[]{[CI] and $J$=2--1 $\\13co$ spectra observed towards sample galaxies. The vertical scale is $T_{\\rm mb}$ in K; the horizontal scale is velocity $V_{\\rm LSR}$ in $\\kms$. For all galaxies, the temperature range in [CI] is four times that in $\\13co$. } \\end{figure*} \\bigskip \\begin{table*} \\caption[]{Line observations log} \\begin{center} \\begin{tabular}{lcccrcccccc} \\hline \\noalign{\\smallskip} Galaxy\t & \\multicolumn{2}{c}{Position} & Adopted & \\multicolumn{3}{c}{[CI] } & \\multicolumn{2}{c}{$J$=2--1$\\13co$ } & \\multicolumn{2}{c}{$J$=4--3 $\\co$} \\\\ & RA(1950) & Dec(1950) & Distance & No. [CI] & Date & $T_{\\rm sys}$ & Date & $T_{\\rm sys}$ & Date & $T_{\\rm sys}$\\\\ & (h m s) & ($^{\\circ}$ $'$ $''$) & (Mpc) & Points &\t & (K) \t & \t & (K) & & (K) \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} NGC~253 & 00:45:05.7 \t& --25:33:38 & 2.5 & 20 & 12/93 & 3770\t & 12/93 & 1695 \t & 11/94 & 9800 \\\\ NGC~278 & 00:49:15.0\t& +47:16:46 & 12 & 1 & 07/96 & 3650\t & 06/95 & 480 \t & 01/01 & 1325 \\\\ NGC~660 & 01:40:21.6\t& +13:23:41 & 13 & 1 & 07/96 & 3065\t & 05/01 &\t350 \t & 08/99 & 3870 \\\\ Maffei~2 & 02:38:08.5\t& +59:23:24 & 2.7 & 1 & 12/93 & 4885\t & 01/96 & 550 \t & 07/96 & 3700 \\\\ NGC~1068 & 02:40:07.2\t& --00:13:30 & 14.4 & 22 & 07/96 & 4000\t & 01/96 & 455\t & 07/96 & 3365 \\\\ IC~342 & 03:41:36.6\t& +67:56:25 & 1.8 & 27 & 11/94 & 4485\t & 02/89 & 1440\t & 04/94 & 2170 \\\\ M~ 82\t & 09:51:43.9\t& +69:55:01 & 3.25 & 6 & 12/93 & 7200\t & 04/93 & 335\t & 10/93 & 9085 \\\\ NGC~3079 & 09:58:35.4\t& +55:55:11 & 18.0 & 7 & 03/94 & 6240\t & 06/95 & 310\t & 03/94 & 5510 \\\\ NGC~3628 & 11:17:41.6\t& +13:51:40 & 6.7 & 8 & 11/94 & 3450\t & 06/95 & 325\t & 03/94 & 2414 \\\\ NGC~4826 & 12:54:17.4\t& +21:57:06 & 5.1 & 2 & 03/97 & 3520\t & 12/93 & 535\t & 12/93 & 2045 \\\\ M~51\t & 13:27:45.3\t& +47:27:25 & 9.6 & 4 & 11/94 & 6600\t & 06/95 & 370\t & 04/96 & 4065 \\\\ M~83\t & 13:34:11.3\t& --29:36:39 & 3.5 & 14 & 12/93 & 4590\t & 06/95 & 430\t & 12/93 & 4360 \\\\ NGC~5713 & 14:37:37.6 & --00:04:34 & 21.0\t& 1 & 02/99 & 8000\t & 12/00 & 515\t & ... & ... \\\\ NGC~6946 & 20:33:48.8\t& +59:58:50 & 5.5 & 17 & 07/96 & 3970\t & 01/96 & 530\t & 07/96 & 2900 \\\\ NGC~7331 & 22:34:46.6\t& +34:09:21 & 14.3 & 1 & 11/96 & 1925\t & 12/97 & 320\t & ... & ... \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{center} \\end{table*} \\begin{figure*}[t] \\unitlength1cm \\begin{minipage}[]{16.4cm} \\resizebox{17.0cm}{!}{\\rotatebox{270}{\\includegraphics*{cifig_1b.ps}}} \\end{minipage} \\end{figure*} \\bigskip \\begin{table*} \\caption[]{[CI], $J$=2--1 $\\13co$, $J$=4--3 $\\co$ intensities} \\begin{center} \\begin{tabular}{lccccccc} \\hline \\noalign{\\smallskip} Galaxy & Offset & \\multicolumn{5}{c}{Center Position}\t\t\t & Area-integrated \\\\ & & $T_{\\rm mb}$([CI]) & $I$([CI]) & $I(\\co 4-3)$ & $I$([CI]) & $I(\\13co)$ & [CI] Luminosity \\\\ & & \\multicolumn{3}{l}{(10$''$)} & \\multicolumn{2}{l}{(22$''$)} & \\\\ & $''$ & (mK) & \\multicolumn{4}{c}{($\\kkms$)} \t \t & ($\\kkms$ kpc$^{2}$) \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} NGC~253\t & 0, 0 & 2615 & 486$\\pm$60 & 1019$\\pm$120 & 290$\\pm$45 & 106$\\pm$13 & 14$\\pm$2.2 \\\\ NGC~278\t & 0, 0 & 100 & 7$\\pm$3 & 9$\\pm$2 & (5$\\pm$1) & 2.6$\\pm$0.4 & (3.1$\\pm$0.7) \\\\ NGC~660\t & 0, 0 & 240 & 50$\\pm$8 & 85$\\pm$12 & (31$\\pm$8) & 7.8$\\pm$1.1 & (26$\\pm$7) \\\\ Maffei~2 & 0, 0 & 190 & 37$\\pm$7 & 405$\\pm$50 & (20$\\pm$7) & 22$\\pm$4 & (0.9$\\pm$0.3) \\\\ NGC~1068 & 0, 0 & 560 & 109$\\pm$12 & 153$\\pm$19 & 49$\\pm$9 & 11$\\pm$2 & 53$\\pm$8.5 \\\\ IC~342\t & 0, 0 & 1030 & 54$\\pm$6 & 209$\\pm$21 & 27$\\pm$7 & 24$\\pm$3 & 1.1$\\pm$0.3 \\\\ M~82 & 0, 0 & 2130 & 224$\\pm$35 & 591$\\pm$95 & 180$\\pm$30 & 60$\\pm$9 & 39$\\pm$6.9 \\\\ NGC~3079 & 0, 0 & 530 & 111$\\pm$18 & 115$\\pm$20 & (70$\\pm$15) & 12$\\pm$3 & 143$\\pm$31 \\\\ NGC~3628 & -17, +5 & 265 & 84$\\pm$11 & 110$\\pm$15 & 38$\\pm$8 & 10$\\pm$2 & 28$\\pm$5.7 \\\\ NGC~4826 & 0, 0 & 135 & 11$\\pm$2 & ... & 17$\\pm$4 & 7.8$\\pm$1.6 & ... \\\\ & -20, +5 & 270 & 86$\\pm$10 & 73$\\pm$9 & (47$\\pm$10) & 15$\\pm$2 & (13$\\pm$3) \\\\ M~51 & 0, 0 & 565 & 28$\\pm$5 & 24$\\pm$4 & (13$\\pm$4) & 8.4$\\pm$1.9 & (8$\\pm$2.7) \\\\ & -12, -12 & 340 & 24$\\pm$5 & ... & ... & ... & ... \\\\ & -12, -24 & 755 & 55$\\pm$9 & ... & ... & ... & ... \\\\ & -24, -24 & 470 & 55$\\pm$9 & ... & ... & ... & ... \\\\ M~83\t & 0, 0 & 685 & 83$\\pm$14 & 270$\\pm$20 & 55$\\pm$8 & 29$\\pm$3 & 3.6$\\pm$0.5 \\\\ NGC~5713 & 0, 0 & $<$90 & 2$\\pm$0.4 & ... & ... & 7.4$\\pm$1.6 & 1.7--2.6 \\\\ NGC~6946 & 0, 0 & 465 & 85$\\pm$9 & 216$\\pm$20 & 44$\\pm$8 & 22$\\pm$3 & 13$\\pm$2.4 \\\\ NGC~7331 & 0, 0 & 30 & 2$\\pm$0.3 &\t ... & ... & 2.5$\\pm$0.6 & 0.8--1.7 \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{center} Note: Offset position (--20,+5) of NGC~4826 is actual nucleus position. \\end{table*} Carbon monoxide (CO), the most common molecule after $\\h2$, is now routinely detected in external galaxies. However, when exposed to energetic radiation, CO is readily photodissociated turning atomic carbon into an important constituent of the interstellar medium. As the ionization potential of neutral carbon is quite close to the dissociation energy of CO, neutral carbon subsequently may be ionized rather easily. As a consequence, [CI] emission primarily arises from interface regions between zones emitting in [CII] and CO respectively (see e.g. Israel et al. 1996; Bolatto et al. 2000). It requires column densities sufficiently high for shielding against ionizing radiation, but not so high that CO selfshielding allows most gas-phase carbon to be bound in molecules. In principle, observations of emission from CO, C$^{\\circ}$ and C$^{+}$ provide significant information on the physical condition of the cloud complexes from which the emission arises (Israel, Tilanus $\\&$ Baas 1998; Gerin $\\&$ Phillips 2000; Israel $\\&$ Baas 2001). Even though the far-infrared continuum and the [CII] line are much more efficient coolants, the various CO and [CI] lines are important coolants for relatively cool, dense molecular gas, contributing about equally to its cooling (Israel et al. 1995; Gerin $\\&$ Phillips 2000). In galaxies, however, studies of the dense interstellar medium are complicated by the effectively very large linear observing beams which incorporate whole ensembles of individual, mutually different clouds. The clumpy nature of the interstellar medium allows UV radiation to penetrate deeply into cloud complexes, so that the CO, [CI] and [CII] emitting volumes appear to coincide when observed with large beamsizes. The physics and structure of such photon-dominated regions (PDR's) has been reviewed most recently by Hollenbach $\\&$ Tielens (1999), whereas their consequent observational parameters have been modelled by e.g. Kaufman et al. (1999). [CII] emission has been observed towards numerous galaxies, both from airborne (the now defunct NASA Kuiper Airborne Observatory) and from spaceborne (the equally defunct Infrared Space Observatory) platforms. In contrast to these [CII] observations, observations of [CI] emission can be performed on the ground, at least in the $\\pci$ transition at 492 GHz. However, atmospheric transparency is poor at such high frequencies and weather conditions need to be unusually favourable for observations of the often weak extragalactic [CI] emission to succeed, even at the excellent high-altitude site of telescopes as the JCMT and the CSO. Consequently, the number of published results is relatively limited. Beyond the Local Group, i.e. at distances larger than 1 Mpc, [CI] has been mapped in bright galaxies such as IC~342 (B\\\"uttgenbach et al. 1992), M~82 (Schilke et al. 1993; White et al. 1994; Stutzki et al. 1997) and NGC~253 (Israel et al. 1995; Harrison et al. 1995), as well as M~83 (Petitpas $\\&$ Wilson, 1998; Israel $\\&$ Baas 2001) and NGC 6946 (Israel $\\&$ Baas 2001). A survey of 13 galaxies, including limited radial mapping of NGC~891 and NGC~6946 was recently published by Gerin $\\&$ Phillips (2000). In this paper, we present a similar [CI] survey of 15 galaxies. We also obtained $J$=2--1 $\\13co$ measurements for all galaxies, and $J$=4--3 measurements for all but two. Taking overlap into account, this survey brings the total number of galaxies outside the Local Group, detected in [CI], to 26. ", + "conclusions": "\\begin{enumerate} \\item We have measured the emission from the 492 GHz line corresponding to the $\\pci$ transition in the centers of fifteen nearby spiral galaxies. In the same galaxies, we have also measured $J = 4-3 \\co$ and $J = 2-1 \\13co$ intensities for comparison with the $\\pci$ line within the framework of radiative transfer models. \\item Rather unlike Galactic sources, the external galaxy centers have [CI] intensities generally exceeding $J$=2--1 $\\13co$ intensities, and in a number of cases approaching $J$=4--3 $\\co$ intensities. \\item The highest area-integrated (i.e. total central) [CI] luminosities are found in the active galaxies NGC~1068 and NGC~3079. Slightly lower luminosities occur in strong starburst nuclei, such as those of NGC~3628 and NGC~6946. Quiescent and weak-starburst nuclei have [CI] luminosities an order of magnitude lower. \\item The observed [CI], $\\13co$ and $\\co$ line ratios, interpreted within the context of radiative transfer models, suggest that the bulk of the observed emission arises in gas with densities $n \\geq 3000 \\cc$ and kinetic temperatures $T_{\\rm kin} \\approx 30 - 60$ K. Depending on the actual density $n$, most galaxy centers should have abundances $N$([CI])/$N$(CO) = 1 -- 3, i.e. [CI] columns just exceeding those of CO. Only relatively quiescent galaxy centers such as those of Maffei~2, IC~342 and NGC~7331 have abundances $N$([CI])/$N$(CO) $\\approx 0.3 - 1.0$ and are dominated by CO just as the comparison starforming regions in the Milky Way and the LMC. \\item The observed [CI] intensities, together with literature [CII] line and far-infrared continuum data, likewise suggest that a significant fraction of the emission originates in medium-density gas ($n = 10^{3}-10^{4} \\cc$), subjected to radiation fields of various strengths ranging from a few times to several thousand times the local Galactic radiation field. \\end{enumerate}" + }, + "0112/astro-ph0112008_arXiv.txt": { + "abstract": "We consider the possibility that the excess of cosmic rays near $\\sim 10^{18}$ eV, reported by the AGASA and SUGAR groups from the direction of the Galactic Center, is caused by a young, very fast pulsar in the high density medium. The pulsar accelerates iron nuclei to energies $\\sim 10^{20}$ eV, as postulated by the Galactic models for the origin of the highest energy cosmic rays. The iron nuclei, after about 1 yr since pulsar formation, leave the supernova envelope without energy losses and diffuse through the dense central region of the Galaxy. Some of them collide with the background matter creating neutrons (from desintegration of Fe), neutrinos and gamma-rays (in inelastic collisions). We suggest that neutrons produced at a specific time after the pulsar formation are responsible for the observed excess of cosmic rays at $\\sim 10^{18}$ eV. From normalization of the calculated neutron flux to the one observed in the cosmic ray excess, we predict the neutrino and gamma-ray fluxes. It has been found that the 1 km$^2$ neutrino detector of the IceCube type should detect from a few up to several events per year from the Galactic Center, depending on the parameters of the considered model. Also future systems of Cherenkov telescopes (CANGAROO III, HESS, VERITAS) should be able to observe 1 - 10 TeV $\\gamma$-rays from the Galactic Center if the pulsar was created inside a huge molecular cloud about $3-10\\times 10^3$ yrs ago. ", + "introduction": "Recently the AGASA collaboration has reported the existence of an extended excess of cosmic rays (CRs) over a narrow energy range $10^{17.9} - 10^{18.3}$ eV from directions close to the Galactic Center (GC) and the Cygnus region with the significance of 4.5$\\sigma$ and $3.9\\sigma$, respectively (Hayashida et al.~1999). The GC excess was confirmed in the analysis of the SUGAR data (Bellido et al.~2001), in which case the observed signal in the energy range $10^{17.9} - 10^{18.5}$ eV is consistent with that from a point-like source of neutral particles, with an estimated flux of $(9 \\pm 3)\\times 10^{-14}$ m$^{-2}$ s$^{-1}$, offset from the true location of the GC by about $7.5^{\\rm o}$. Hayashida et al.~(1999) suggested that such an anisotropy of CRs in a narrow energy range can be explained naturally by neutrons due to the fact that particles with such energies are able to reach the Earth before decaying from distance of the GC. Numerical simulations of charged particles (protons) propagation in reasonable models of the Galactic magnetic field give rise to the more extended source which may even be significantly shifted in the sky from the original direction towards the source (e.g. Clay~2000, Bednarek, Giller \\& Zieli\\'nska~2001). However, the results of charged hadrons propagation strongly depend on the details of the magnetic field model which is not well known at present. The production of neutrons in discrete galactic sources was already discussed some time ago by Jones~(1990), Sommers \\& Elbert~(1990) and Bednarek~(1992) in relation to reports of the excess of EeV particles from the direction of Cygnus X-3 (Cassiday et al.~1989, Teshima et al.~1990). It has recently been suggested that in the GC region neutrons might be produced in collisions of hadrons with matter (Takahashi \\& Nagataki~2001). Hadrons can be accelerated by a massive black hole associated with the Sgr A$^*$ (Levinson \\& Boldt~2000), or by the shock waves of supernovae which explode into their own stellar winds (Rhode, Enslin \\& Biermann~1998). The Galactic Center region (inner $\\sim 50$ pc) is rich in many massive stellar clusters with a few to more than 100 OB stars (Morris \\& Serabyn~1996, Blum et al.~2001). These stars should soon explode as supernovae. In fact recent multiple supernova explosions in the GC region ($10^3$ supernovae in the past $10^5$ years) are suggested by the observations of the diffuse hot plasma emitting X-rays (Yamauchi et al.~1990). For example one remnant of such a young supernova with the age of $\\sim 80$ years (G0.570-0.018) has recently been reported by Senda et al.~(2001). Since it is expected that pulsars are formed in explosions of such massive stars, we can expect that the GC region should contain some young pulsars, a number of them being $10^2 - 10^3$ yrs old. Motivated by these observational results we may assume that at least one of these young pulsars, formed in a supernova explosion of the type Ib/c, has parameters which allow the acceleration of iron nuclei to energies $\\sim 10^{20}$ eV, as postulated by some models for the Galactic origin of the highest energy cosmic rays (Blasi, Epstein \\& Olinto~2000; De Goubeia Dal Pino \\& Lazarian~2000). We suggest that neutrons from desintegration of iron nuclei, which are accelerated by such energetic pulsar in the GC region, can be responsible for the observed excess of the cosmic rays with energies $\\sim 10^{18}$ eV. Note that the iron nuclei with energies $\\sim 10^{20}$ eV and neutrons with energies $\\sim 2\\times 10^{18}$ eV have the same Lorentz factors. In order to test this hypothesis we predict the neutrino and $\\gamma$-ray fluxes accompaning the process of neutron injection by iron nuclei. ", + "conclusions": "The SUGAR group estimates the flux of particles which causes reported excess of the cosmic ray particles in the energy range $10^{17.9}-10^{18.5}$ eV on $(9\\pm 3)\\times 10^{-14}$ m$^{-3}$ s$^{-1}$ (Bellido et al.~2001). If this excess is caused by neutrons produced in the pulsar model discussed here, then the expected flux of neutrons can be compared with the observed one. Basing on this normalization we predict the fluxes of neutrinos and gamma-rays on Earth. This procedure allows us to derive the free parameter of our model (i.e. the efficiency $\\xi$ of iron acceleration by the pulsar) and limit the age of the pulsar for other fixed parameters, $P, B, R_{\\rm c}, n_{\\rm c}, B_{\\rm c}$, which are in fact constrained by the observations. We consider five different sets of parameters describing our scenario: model (I) R = 10 pc, $n = 10^3$ cm$^{-3}$, $B_{\\rm c} = 10^{-4}$ G, $t_{\\rm obs} = 10^4$ yr; (II) $t_{\\rm obs} = 3\\times 10^3$ yr and other parameters as above; (III) $t_{\\rm obs} = 10^3$ yr and other parameters as above; (IV) R = 50 pc, $n = 10^2$ cm$^{-3}$, $B_{\\rm c} = 3\\times 10^{-5}$ G, $t_{\\rm obs} = 10^4$ yr; and (V) $t_{\\rm obs} = 3\\times 10^3$ yr and other parameters as in (IV). They all concern two sets of parameters for the medium in which the pulsar is formed, and differ in the pulsar's age which is not constrained by any observations. In all these models we assume that the pulsar is born with $B = 6\\times 10^{12}$ G and $P_{\\rm o} = 2$ ms. Normalizing the predicted neutron flux to the observed excess of CR particles we derive the value of the parameter $\\xi$ which has to be $\\xi\\approx 1$ (model I), $0.18$ (II), $0.03$ (III), $0.3$ (IV), and $0.09$ (V). \\begin{table} \\caption{Gamma-rays and neutrinos from the Galactic Center.} \\begin{tabular}{|c|c|c|c|c|} \\hline Model & $N_\\gamma$($>1$ TeV)& $N_\\gamma$($>10$ TeV) & $N_\\nu^{\\rm a}$ & $N_\\nu^{\\rm na}$ \\\\ \\hline (I) & $4.3\\times 10^{-12}$ & $2.2\\times 10^{-12}$ & 23 & 30 \\\\ \\hline (II) & $8.7\\times 10^{-13}$ & $6.6\\times 10^{-13}$ & 11 & 16 \\\\ \\hline (III) & $1.4\\times 10^{-13}$ & $1.25\\times 10^{-13}$ & 4.2 & 7.1 \\\\ \\hline (IV) & $2.5\\times 10^{-13}$ & $1.7\\times 10^{-13}$ & 5.3 & 8.8 \\\\ \\hline (V) & $6.7\\times 10^{-14}$ & $5.5\\times 10^{-14}$ & 2.0 & 3.8 \\\\ \\hline \\end{tabular} \\label{table1} \\end{table} Using the above estimates for $\\xi$ we can now predict the expected fluxes of $\\gamma$-rays and muon neutrinos and antineutrinos in the case of every model. The integral spectra of $\\gamma$-rays from the GC region are presented in Fig.~\\ref{fig3}, together with the sensitivities of the present HEGRA telescope system and the planned next generation of telescopes, i.e. CANGAROO III, HESS, VERITAS. We also report in Table~\\ref{table1} the $\\gamma$-ray fluxes above 1 TeV and 10 TeV in units cm$^{-2}$ s$^{-1}$. Although the $\\gamma$-ray spectra have a maximum above 10 TeV for all models, the $\\gamma$-ray fluxes in the energy range 1-10 TeV produced in models, (I) $\\sim 2.1\\times 10^{-12}$ cm$^{-2}$ s$^{-1}$, and (II) $\\sim 2.1\\times 10^{-13}$ cm$^{-2}$ s$^{-1}$, and probably also in (IV) $\\sim 8\\times 10^{-14}$ cm$^{-2}$ s$^{-1}$, should be observed by the future systems of Cherenkov telescopes of the CANGAROO III, HESS, and VERITAS type. Models (III) and (V) predict fluxes below the sensitivity limit of these Observatories. Only the HEGRA Collaboration observed the Galactic disk including the GC region (P\\\"ohlhofer et al.~1999). The upper limit on the possible sources, equals 1/4 Crab in the Galactic plane which is above the $\\gamma$-ray flux predicted even by the model (I). However, since the GC region can be observed by this experiment only at zenith angles larger than $60^{\\rm o}$, this limit does not refer to the GC region. In Fig.~\\ref{fig4} we show the muon neutrino and anti-neutrino spectra expected in the discussed model. At energies $> 10$ TeV, these spectra are above the expected flux of atmospheric neutrino background (ANB) and also above the 3 yr sensitivity limit of the planned large size neutrino detector IceCube (Hill~2001). We estimate the number of muon neutrino events during one year in the IceCube detector basing on the calculations of the likelihood of detecting such neutrinos by a detector with a surface area of 1 km$^2$ obtained by Gaisser \\& Grillo~(1987). The results of these calculations are shown in Table~1. We distinguish the case of neutrinos coming to the neutrino detector from directions close to the horizon, i.e. not absorbed by the Earth ($N_{\\nu}^{\\rm na}$), and neutrinos which arrive moving upward from the nadir direction, i.e. partially absorbed by the Earth ($N_{\\nu}^{\\rm a}$) (for absorption coefficients see Gandhi~2000). From Table~1 it is clear that the IceCube detector should detect a few up to several neutrinos per year from the Galactic Center region provided that the excess of cosmic rays at $\\sim 10^{18}$ eV from the GC region is caused by neutrons from desintegrations of iron nuclei, accelerated by a very fast pulsar. The detection of the predicted fluxes of neutrinos from the Galactic Center (or lack) will also place constraints on the recent model of extremely high energy cosmic ray production in the pulsar scenario (Blasi, Epstein \\& Olinto~2000), since the parent iron nuclei which inject neutrons with energies $\\sim 10^{18}$ eV, have to be accelerated to energies $\\sim 10^{20}$ eV. \\begin{figure} \\vspace{6.truecm} \\special{psfile=mb920fig4.eps voffset=0 hoffset=-10 hscale=55 vscale=50} \\caption{The differential spectra of muon neutrinos and antineutrinos produced in hadronic interactions of iron nuclei with the matter of a molecular cloud at the Galactic Center for the models discussed in Fig.~3. The dashed curves indicate the atmospheric neutrino background, ANB, (Lipari~1993) within a $1^{\\rm o}$ of the source and the dotted line shows the 3 yr sensitivity of the IceCube detector (Hill~2001).} \\label{fig4} \\end{figure} The EGRET detector on board the Compton GRO has detected a strong $\\gamma$-ray source with luminosity of $\\sim 2\\times 10^{37}$ erg s$^{-1}$ (Mayer-Hasselwander et al.~1998). This emission seems to come from the extended region with the radius of $\\sim 80$ pc around the GC. In another paper (in preparation) we suggest that this emission can be explained in terms of the general scenario discussed here if a part of energy of the accelerated iron nuclei is transferred to the relativistic positrons due to the resonant scattering in the pulsar shock region (see Gallant \\& Arons~1994). These positrons, accumulated in the cloud, produce high energy radiation in synchrotron and inverse Compton processes." + }, + "0112/astro-ph0112522_arXiv.txt": { + "abstract": "We explore some of the consequences of Dark-Matter--photon interactions on structure formation, focusing on the evolution of cosmological perturbations and performing both an analytical and a numerical study. We compute the cosmic microwave background anisotropies and matter power spectrum in this class of models. We find, as the main result, that when Dark Matter and photons are coupled, Dark Matter perturbations can experience a new damping regime in addition to the usual collisional Silk damping effect. Such Dark Matter particles (having quite large photon interactions) behave like Cold Dark Matter or Warm Dark Matter as far as the cosmic microwave background anisotropies or matter power spectrum are concerned, respectively. These Dark-Matter--photon interactions leave specific imprints at sufficiently small scales on both of these two spectra, which may allow us to put new constraints on the acceptable photon--Dark-Matter interactions. Under the conservative assumption that the abundance of $10^{12} M_\\odot$ galaxies is correctly given by the Cold Dark Matter, and without any knowledge of the abundance of smaller objects, we obtain the limit on the ratio of the Dark-Matter--photon cross section to the Dark Matter mass $\\frac{\\sigma_\\SGAMMADM}{m_\\DM} \\lesssim 10^{-6} \\frac{\\sigma_\\THM}{100 \\UUNIT{GeV}{}} \\simeq \\EE{6}{-33} \\UUNIT{cm}{2} \\UUNIT{GeV}{-1}$. \\vspace*{0.5cm} \\noindent {\\footnotesize Preprint numbers: SPhT-Saclay T01/147, {\\tt astro-ph/0112522}} ", + "introduction": "The nature of Dark Matter particles remains one of the major challenges for both fundamental physics and astrophysics. Whereas Cold Dark Matter (CDM) perfectly explains the formation of large scale structure~\\cite{Peacock:1999ye} on scales greater than $1 \\UUNIT{Mpc}{}$, there seem to be various discrepancies on smaller (subgalactic) scales. Some of these come from the following. \\begin{enumerate} \\item \\label{prob1} $N$-body CDM simulations, which give cuspy halos with divergent profiles toward the center~\\cite{NFW}, in potential disagreement with the galaxy rotation curves~\\cite{binney} and with observations from gravitational lensing~\\cite{flores}; \\item \\label{prob2} Bar stability in high surface brightness spiral galaxies which also demands low-density cores~\\cite{debattista}; \\item \\label{prob3} CDM models which for years have been seen to yield an excess of small scale structures~\\cite{ss1985}. Numerical simulations~\\cite{Klypin:1999uc} found $1$--$2$ orders of magnitude more satellite galaxies than what is observed~\\cite{Moore:1999wf}. However, recent work~\\cite{stoehr,haya,chiu} indicate that there may be no problem for the galaxy mass function after all. \\item \\label{prob4} The formation of disk galaxy angular momentum, which is much too small in galaxy simulations~\\cite{NFW}. \\end{enumerate} Problems~\\ref{prob3} and~\\ref{prob4} can be solved with the usual Warm Dark Matter (WDM) which experiences free streaming and hence suppresses power on small scales~\\cite{ss1988,sommerlarsen,nbody}. Such a particle physics candidate is easy to find in a minimalistic extension of the Standard Model, namely, a sterile neutrino~\\cite{dodelson,Shi:1999km,Dolgov:2000ew,Abazajian:2001nj}. However, collisionless WDM does not solve problems~\\ref{prob1} and~\\ref{prob2} (see Ref.~\\cite{knebe}) and one is therefore forced to propose more complicated models like scenarios of nonthermal production of Weakly Interacting Massive Particles (WIMPs)~\\cite{zhang} for instance. On the other hand, Strongly Interacting Dark Matter (SIDM) has been suggested~\\cite{spergel} to solve problems~\\ref{prob1} and~\\ref{prob2} and does so successfully provided the cross section is within the range $\\EE{2}{-25} \\UUNIT{cm}{2} \\UUNIT{GeV}{-1} \\lesssim \\sigma_\\SDMDM / m_{\\DM} \\lesssim \\UEE{-23} \\UUNIT{cm}{2} \\UUNIT{GeV}{-1}$~\\cite{daveS,wandelt} (where $m_\\DM$ and $\\sigma_\\SDMDM$ are the Dark Matter mass and self-interaction cross section respectively). Problem~\\ref{prob3} is also partly solved in this scenario~\\cite{daveS} but the survival of galactic halos exclude the range $\\EE{6}{-25} \\UUNIT{cm}{2} \\UUNIT{GeV}{-1} \\lesssim \\sigma_\\SDMDM / m_{\\DM} \\lesssim \\EE{2}{-20} \\UUNIT{cm}{2} \\UUNIT{GeV}{-1}$~\\cite{GnedinO}. Furthermore, since the inner regions of massive clusters are elliptical, one must have $\\sigma_\\SDMDM / m_\\DM \\lesssim \\EE{3}{-26} \\UUNIT{cm}{2} \\UUNIT{GeV}{-1}$~\\cite{miralda}. One therefore concludes that the allowed cross section is slightly too small and that SIDM, on its own, cannot solve problems~\\ref{prob1}--\\ref{prob4}. This paper is finally motivated by the recent findings~\\cite{bfs2001} that either Dark-Matter--photon or Dark-Matter--neutrino interactions can transfer to Dark Matter the damping that the photon or neutrino fluids undergo. This process, characterized in the simplest cases by an exponential cutoff in the matter power spectrum, was referred to as ``induced damping'' in Ref.~\\cite{bfs2001}. In particular, by requiring that the damping induced by relativistic particles does not wash out the Dark Matter primordial fluctuations responsible for the formation of the smallest galaxies, it was found that the ratio of the corresponding cross sections to the Dark Matter mass must satisfy $\\sigma_\\SGAMMADM / m_\\DM < \\UEE{-30} \\UUNIT{cm}{2} \\UUNIT{GeV}{-1}$ and $\\sigma_\\SNUDM / m_\\DM < \\UEE{-34} \\UUNIT{cm}{2} \\UUNIT{GeV}{-1}$. It was then suggested that, at the edge to be satisfied, these constraints could provide an alternative scenario for Warm Dark Matter. However, the exact value of these cross sections as well as the shape of the resulting power spectrum depend on the details of the interactions history of the fluids. We therefore determine---in this paper---the transfer function resulting from non-negligible interactions between Dark Matter and photons. The case of neutrino--Dark-Matter interactions will be examined in a subsequent paper. The effects thereof are naturally different from those due to self-interactions. However, one should keep in mind that realistic a Dark Matter particle probably has interactions both with itself and with other particles. In Sec.~\\ref{celine_part}, we discuss the motivations for such Dark Matter particles. In Sec.~\\ref{alain_part} we describe the effect of Interacting Dark Matter\\footnote{In the following, we shall adopt the notation of Interacting Dark Matter for convenience but the reader has to remember that only Dark-Matter--photon interactions are investigated in this paper.} (IDM) on the early evolution of cosmological perturbations and, in Sec.~\\ref{steen_part}, we give an analytical fit to the main features produced by IDM on the matter power spectrum. In the Conclusion, we discuss the main result of our work, namely that, for Dark Matter coupled to photons, the original fluctuations are damped as soon as they enter the horizon, because of the impeded growth due to this coupling. We also discuss a few implications of this result. ", + "conclusions": "We have considered the effect of Dark-Matter--photon interactions on the evolution of primordial Dark Matter fluctuations. Rather than growing when they enter the horizon, the fluctuations stay of constant amplitude, as do the photon fluctuations. This impeded growth appears as a damping compared to the original amplitude of the fluctuations. Moreover, fluctuations on scales much below the size of the horizon are seen to couple to the photons at a rate much lower than the rate at which they oscillate. This is an additional, new, damping process we have called weak coupling. As a result, horizon-size Dark Matter fluctuations are seen to be damped if coupled to photons. The usual (exponential) damping sets in for much smaller scales. We hence have obtained a new constraint on the allowed cross sections. This bound~(\\ref{wcbound}) is two orders of magnitude stronger than the necessary condition obtained by considering the exponential damping of Dark Matter fluctuations induced by the photon interactions~\\cite{bfs2001}. This new bound reads \\begin{equation} \\label{wcbound} \\sigma_\\SGAMMADM / m_\\DM \\lesssim 10^{-32} \\UUNIT{cm}{2} \\UUNIT{GeV}{-1} . \\end{equation} The maximum allowed value reduces the matter power spectrum in a way corresponding to a conventional Warm Dark Matter particle with a mass of about $1 \\UUNIT{keV}{}$ but leaves the cosmic microwave background anisotropies undisturbed. Using recent bounds on the scale of reduction of the matter power spectrum thus allows us to put new bounds on the allowed photon--Dark-Matter cross section. This corresponds to a universe which is well transparent to photons: the free mean path of a photon due to the interactions with Dark Matter in a halo core of mass density $0.02 M_\\odot \\UUNIT{pc}{- 3}$ is of order $\\EE{6}{4} \\UUNIT{Gpc}{}$, and the optical thickness toward the (usual) last scattering surface is below $10^{- 5}$. The value~(\\ref{wcbound}) nevertheless remains quite large compared to the theoretical estimates usually encountered for weakly interacting particles, although there are no compelling reasons to exclude it. Anyway, this leaves open new possibilities as far as the nature of Dark Matter is concerned. The lower bound~(\\ref{wcbound}) implies that Dark Matter decouples from the photons before the collisional Silk damping is at work, leaving an oscillating, power-law, damped matter power spectrum. This new damping regime bears some similarities with the free streaming case, although here the Dark Matter and photon fluids are undoubtedly coupled. Such Dark Matter particles therefore appear to be good Warm Dark Matter candidates, with features in the matter power spectrum different from the conventional WDM at very small scales." + }, + "0112/astro-ph0112244_arXiv.txt": { + "abstract": "Rapid growth of neutrino-nucleon cross-sections at high energies due to hypothetical hard pomeron enhancement of \\nuN-structure functions is discussed. Differential and integral hadron moments, which, together with cross-sections, define rates of hadron-electromagnetic cascades in a neutrino detector are calculated for different power-law decreasing neutrino spectra. For comparison two small-$x$ extrapolation schemes are discussed. First includes Regge theory inspired hard pomeron enhancement of \\nuN-structure functions. The second is obtained with the help of trivial extrapolation of perturbative \\QCD\\ structure functions from the large $x$ region to the small $x$ one. Implications of hard pomeron effects for cross-sections and hadron moments are demonstrated. The most pronounced manifestations are found in integral hadron moments for the case of charged current electron (anti)neutrinos scattering off nucleons. ", + "introduction": "This paper continues the discussion of some specific features of \\nuN\\ \\emph{Deep Inelastic Scattering (DIS)} at extremely high, up to $E_\\nu \\sim 1\\times 10^{12}$~GeV, energies that was started at the previous Seminar \\emph{'NPCS-2000'} \\cite{GYNPQS}. The main idea of our approach is to extend the successful small-$x$ description of $F_2^{ep}(x,Q^2)$ \\emph{Structure Function (SF)} by A. Donnachie and P. V. Landshoff \\emph{(DL)} \\cite{DL} to \\nuN-\\SFs, namely, to $F_2^{\\nu N}(x,Q^2)$, $F_3^{\\nu N}(x,Q^2)$. \\DL\\ claim that record small-$x$ $ep$-scattering data by \\emph{HERA} \\cite{HERA} may be successfully explained with the help of a simple combination of several Regge theory inspired non-perturbative pomerons. The most important are \\emph{'soft'} pomeron (with intercept $\\sim 1.08$) and \\emph{'hard'} pomeron (with intercept $\\approx 1.4$). First prevails at small $Q^2$, while the latter dominates at large $Q^2$. Moreover, \\DL\\ argue that perturbative \\QCD\\ (\\pQCD) fails at small $x < 10^{-5}$ and that its validity at $x \\sim 10^{-4}\\div 10^{-5}$ is a pure fluke. However, one should keep in mind that \\DL's approach neglects the other, non-leading, poles and cuts in the complex angular momentum $l$-plane. This common feature of pomeron physics causes this model to violate the unitarity at $E_\\nu \\rightarrow \\infty$. Using a nontrivial generalization of \\DL's $F_2^{ep}(x,Q^2)$ \\SF\\ description to the \\nuN-scattering case, we have constructed $F_2^{\\nu N}(x,Q^2)$ and $F_3^{\\nu N}(x,Q^2)$ \\SFs, presumably valid in the whole range of kinematic variables $0 \\leq x \\leq 1$ and $0 \\leq Q^2 \\leq \\infty$ \\cite{GYNPQS,GY}. At $x \\gtrsim 10^{-5}$ they are chosen to coincide with \\pQCD\\ parameterization by \\emph{CTEQ5} collaboration \\cite{CTEQ}, while in the small-$x$ region these \\SFs\\ are driven by the analogous Regge theory inspired description. A special interpolation procedure, developed in Ref.~\\cite{GY}, allows to meet smoothly these different, both over $x$ and $Q^2$, descriptions of \\SFs\\ at low and high $x$. In Ref.~\\cite{GYNPQS} these \\SFs\\ were denoted by \\emph{DL+CTEQ5}, indicating that they have their origin in the interpolation between \\DL\\ and \\pQCD\\ descriptions. In parallel there were considered \\SFs\\ obtained via simple extrapolation of \\pQCD\\ \\SFs\\ from $x \\ge 1\\times 10^{-5}$ to the small-$x$ region: \\begin{equation} \\label{log} F_i^{\\nu N,Log+CTEQ5}(x 10^6$ the time required is measured in years, even on the most powerful supercomputers. A number of authors have suggested ways around this limitation, see {\\it e.g.} Oh Spergel \\& Hinshaw 1999, Hivon, \\etal \\ 2001 and Dor\\'{e}, Knox \\& Peel 2001. Such methods usually rely on symmetries in the data, {\\it i..e}, axial symmetry of the noise, while neural networks need to make no assumptions about symmetries. Neural networks provide an alternative for astrophysical parameter estimation.They have been used previously in astronomy for galaxy classification \\cite{Lahav96,Andreon00} and periodicity analysis of unevenly sampled data as applied to stellar light curves \\cite{Tagliaferri00}. They have also been used to analyze stellar spectra \\cite{BailerJones97,BailerJones98,BailerJones00}, with results comparable to traditional methods. However, Bailer-Jones {\\it et. al.} compared data to a deterministic model (stellar spectra), whereas cosmological applications examine random patterns drawn from parameterized stochastic models. We demonstrate the generality of our algorithm by considering different problems with the same network architecture. We find the computational cost for training the network, in the context of CMB anisotropy, requires $O(N^{3/2})$ operations and thus provides a substantial improvement over brute-force maximum-likelihood methods. The stochastic nature of our models is fundamental: their starting point is the quantum nature of the early universe. What the models predict are the parent populations from which individual realizations are to be drawn. Our single observed universe is assumed to be such a realization. The question becomes: given a particular model, which parameters that define the parent population best describe the observed data? This is to be contrasted with the problems of steller spectra or galaxy classification, which have at their heart either a deterministic model or a pre-assigned catalogue of types. This stochastic nature means any analysis of the observed data must rely on some statistical test. Leaving aside for the moment the computational challenges of traditional maximum likelihood methods, this raises the problem of knowing the best statistic for distinguishing between different parameters or models (we can view model selection as a discrete parameter to be estimated). Maximum likelihood methods require an {\\it a priori} definition of a goodness-of-fit function. The choice of a goodness-of-fit function is not always obvious and is particularly acute for 2D and 3D surveys. Much of the information lies in the {\\it phase} features of these surveys. Statistical tests can fail badly in detecting phase features, as witness the large literature devoted to the relatively simple problem of edge detection in 2D data sets (see, e.g., Hough\\ 1962 and Davis\\ 1975). Topological tests such as the genus or other Minkowski functionals have been applied to 2- and 3-D maps, \\cite{Gott90,Kogut96} but the relative power of these statistics is poor \\cite{PhillipsKogut01}. Neural networks, in contrast, do not require specification of a single statistic of {\\it a priori} interest. As the network is trained, it determines how it will discriminate between competing models. All observations add instrumental effects to the desired physical signal. Any analysis of the data must model these effects. For neural networks, this present no great challenge as long as we can model the effects. For either irregularly sampled time series data or 2D/3D survey data with missing patches, any analysis based on Fourier methods will suffer from alaising of power. With neural networks, this problem does not arise; by excluding from the simulations the data points missing in the observed data, the networks will learn the effects of the gaps. The inclusion of noise to the simulations may increase the number of training passes needed, but will not prevent the use of neural networks for parameter estimation. ", + "conclusions": "We have shown neural networks can be used as a tool for astrophysical parameter estimation. For specificity, we have worked in the cosmological context of CMB anisotropy maps where the stochastic nature of the problem is fundamental. The results are insensitive to noise levels and sampling schemes typical of large astrophysical data sets and provide parameter estimation comparable to maximum likelihood techniques. If we classify parameter estimation techniques as to whether they are forward or reverse algorithms, we see the real strength of neural networks. Maximum-likelihood methods are an example of reverse algorithms. They start with the statistic under consideration and work backwards, inverting a covariance matrix, to the likelihood function used to compare different parameter choices. Forward algorithms provide a way to avoid the high computational costs of inverse methods. Typically, it is much simpler to generate model predictions at each sampled point in parameter space than to compute the matrix inverse and determinant required for maximum likelihood techniques. Forward algorithms trade many realizations of synthetic data sets computed at specific parameter values for the computationally infeasible matrix inversion. Neural networks are such an algorithm; synthetic data sets are used to both train and sample the networks. This gives us our speed improvement. Since either maximum likelihood or neural networks can be viewed as the ``machinery'' for parameter estimation, the fundamental information flow stays the same (see Figure \\ref{fitting-model}). The statistical confidence levels for the fitted parameters are always accessible. When the ``machinery'' is sampled with independent synthetic data, we can determine the probabilities for making correct or incorrect parameter identifications. Such sampling also gives us direct access to the statistical power \\cite{PhillipsKogut01}. While training, the information distinguishing the different parameters is encoded in the weights. Interperting the resulting weight matrices is not usually possible (as compared to the Fisher matrix, Eqn \\ref{fisher_def}). Using independent sampling of the network to derive the probability distributions needed for, {\\it e.g.} Bayesian analysis, means we do not need direct access to the information in the weight matrices. Neural networks do not require that we specify one single statistic of {\\it a priori} interest, a limitation of maximum likelihood As the network is trained, it determines how it will discriminate. The information required to separate different parameter points comes from the training set simulations. This lack of a need for a goodness-of-fit function can make neural networks ideal for questions that have been traditionally hard to to answer. These include probing the global topology of the universe \\cite{Lachieze-Rey95,Levin98} and exploring the viable of non-Gaussian models of structure formation. Both of these problems involve detecting global features in the data sets and there is no strong consensus as to the best statistic to use. The often used power spectrum fails to capture enough information to decisively test either hypothesis. If there is a non-trivial topology to the universe, then the isotropy of the universe assumed when using the power spectrum is broken. Neural networks may turn out to be the ideal method for detecting the global features that would be present if indeed the universe does have a non-trivial topology. Any non-Gaussianity present in the primordial fluctuations that seed structure formation is beyond the power spectrum's ability to identify, since it measures the second moment. The bispectrum \\cite{Heavens98,Ferreira98,PhillipsKogut01} (the harmonic conjugate of the three point correlation function) is often used to detect evidence of a departure from pure Gaussian. Neural networks, when combined with a non-Gaussian theory, will provide a complement to bispectrum tests. Along with the explicit example we have presented here of Cosmic Microwave Background anisotropy data, neural networks as a tool for parameter estimation has application to other types of data. Redshift surveys such as the 2-Degree Field and the Sloan Digital Sky Survey measure the redshift and position on the sky of a large number of galaxies ($N \\sim 10^6$), sampling the quasi-linear regime $~\\sim 100 h^{-1}$ Mpc where $h$ is the Hubble constant in units 100 km s${}^{-1}$ Mpc${}^{-1}$. The observed redshift is the sum of the Hubble flow and the peculiar velocity induced by gravitational acceleration in the evolving density field. Coherent flows on large scales produce artifacts in the redshift distribution compared to real space. Galaxies on the far side of an overdensity tend to flow toward the center (hence toward the observer) so that their peculiar velocities subtract from the Hubble flow, making them appear closer than they really are. Galaxies on the near side move the opposite direction, so their peculiar velocities add to the Hubble flow. The net result is an apparent enhancement in the galaxy density in redshift space on scales of superclusters, compressing the region along the line of sight to the observer. The amplitude of this ``bull's-eye'' effect depends on the matter density $\\Omega_m$ on scales comparable to superclusters of galaxies and can be used to determine $\\Omega_m$ in model-independent fashion \\cite{Praton97,Melott98}. Estimating $\\Omega_m$ from distortions in redshift space has several problems in practice. The first is defining a statistic to quantify the bull's-eye enhancement in concentric rings about the origin. \\cite{Melott98} use a large number of simulations to develop an empirical statistic defined as the ratio of rms spacing between upcrossings in isodensity contours in the redshift (radial) direction to that in the orthogonal (azimuthal) direction. It is thus a local statistic in that it compares high-density regions only to other nearby regions, and operates only on a single slice of redshift space after smoothing and contouring. Neural nets, by contrast, offer a {\\it global} test by comparing each region of the density field to all other regions simultaneously, and can easily be extended across the entire three-dimensional survey. No {\\it a priori} statistic need be identified, nor do neural nets require contouring of the density field, thus avoiding the need to ``fine-tune' the selection of contour levels. Neural networks offer a promising approach to cosmological parameter estimation, where the statistical properties of the primordial matter and energy distribution provide one of the few falsifiable tests of the standard inflationary paradigm. They do this at a computational cost much lower than traditional maximum likelihood methods. In the context of CMB anisotropy maps, this cost, $O(\\Ninput^{1.5})$, is better or comparable to the best approximate methods." + }, + "0112/astro-ph0112503_arXiv.txt": { + "abstract": "In the local part of the universe, a majority of the total mass is in only a small number of groups with short dynamical times. The Virgo Cluster dominates far out of proportion to its contribution in light. Ninety percent of the total mass is in groups with $12.5 < {\\rm log} M/M_{\\odot} < 15$. ", + "introduction": "It would hardly be surprising if the relationship between mass and light varies with environment, even if structures form from an invariant baryon to dark matter fractionation. Dark matter could be more dispersed than the baryons that are manifested in stars or detectable gas. This possible differentiation would create a `bias' between what is seen and what exists [1]. Recent modelling [2,3] suggests there could be a complex relationship betwen mass and light, that dark matter may be underrepresented by light at both extremes of low density and high density. Mass-to-light may grow with scale around galaxies to an asymptotic limit [4]. This paper presents observational evidence from the motions of galaxies that there are mass-to-light differences with environment. ", + "conclusions": "Figure~3 shows the inventory of the clustered mass in the local region. Mass contributions are summed over all groups and individual galaxies within a distance of $25 h_{75}^{-1}$~Mpc and with $\\vert b \\vert > 30^{\\circ}$ (distances based on a numerical action kinematic model). In the case of groups, masses come from application of the virial theorem. In the instances of pairs or triples where the observed dispersion in velocities is dominated by measurement uncertainties, or in the case of single galaxies, masses are inferred assuming $M/L_B = 100 M_{\\odot}/L_{\\odot}$, a round off of the mean result for groups represented by the dotted lines in Fig 2. The volume contains a total luminosity of $1.0 \\times 10^{13} L_{\\odot}$ in cataloged galaxies and a total mass of $2.1 \\times 10^{15} M_{\\odot}$ (only 10\\% of this mass total is inferred from the $M/L$ assumption). Within this nearby volume, there should be reasonable completion down to $0.1 L_B^{\\star}$. The overall $M/L \\sim 200 M_{\\odot}/L_{\\odot}$ is consistent with $\\Omega_m \\sim 0.2$ [9]. The mean density in this local region is $4.4 \\times 10^{-30}h_{75}^2$ gm/cm$^3$, indicating this region has twice the cosmic mean density if $\\Omega_m = 0.2$. Hence our census accounts for all the mass anticipated by the numerical action dynamical modelling. There are two striking points to note with Fig 3. One point is that 90\\% of the mass is in bound entities with ${\\rm log}M > 12.5$. The other point is that the Virgo Cluster, with 15\\% of the light in this volume, has over 40\\% of the mass! Of course, the statistics above ${\\rm log}M = 14$ are inadequate with this local sample (Fornax Cluster is the second massive entity). The information provided by groups about the distribution of dark matter on scales of hundreds of kiloparsecs and the information provided by galaxy flows about dark matter on scales up to tens of megaparsecs give rise to a consistent picture. $M/L_B$ values in dynamically evolved regions can be an order of magnitude higher than in the great majority of places that are dynamically young. The evidence from both galaxy motions and an inventory of groups suggests $M/L_B\\sim 200~M_{\\odot}/L_{\\odot}$ overall, consistent with $\\Omega_{matter}\\sim 0.2$. Similar results are found from wide field weak lensing studies [11]. Within the high latitude, inner 25 Mpc region of the group sample, roughly 60\\% of the mass is in 14 low crossing time groups dominated by early galaxy types that contribute a quarter of the light. We seem to live in a curious universe where a majority of the clumped matter is in the modest percentage of locations with crossing times $<2$ Gyr. Groups containing familiar luminous galaxies all have masses above $\\sim 10^{11}~M_{\\odot}$. Candidate groups of only dwarf galaxies are identified with masses near this $10^{11}~M_{\\odot}$ limit. These associations are probably bound, whence they would contain mostly dark matter. It is reasonable to speculate that there are low mass halos without any stars or gas. However most of the mass of the universe seems to be in collapsed regions with $10^{12.5}-10^{15}~M_{\\odot}$." + }, + "0112/astro-ph0112029_arXiv.txt": { + "abstract": "{ IRS~13E is an infrared, mm and X-ray source in the Galactic Centre. We present the first {\\it Chandra} X-ray spectrum for IRS~13E and show that it is consistent with a luminous and highly absorbed X-ray binary system. Since the X-ray luminosity is too large for a solitary star, our interpretation is that of an early-type long-period binary with strong colliding winds emission. This naturally explains the observed X-ray spectrum and count rate as well as its lack of significant short term variability. Due to the short lifetime of any nebula $0.2$~pc from the putative central super-massive black-hole, we argue that the primary of IRS~13E has exited the LBV phase in the last few thousand years. ", + "introduction": "It is probable that Sgr A*, the compact, nonthermal radio source at the Galactic Centre (GC) is a 2--3$\\times10^6 M_{\\sun}$ black-hole \\citep[for a recent review see][]{MF01}. Pervading the central parsec of the Milky Way is a cluster of a few dozen HeI and early-type stars \\citep{SMBH90,GTKKT96}. One of these stars, IRS~13E, has been identified as a Wolf-Rayet (WR) star with spectral class WN10 \\citep{Najarroetal97} and lies within IRS~13, a compact HII region. The IRS~13 complex, dominated by IRS~13E, has been identified as a HeI, Pa-$\\alpha$, [FeIII], and HeII line source \\citep{LKG93, LPFA95, Krabbeetal95, Stolovyetal99}. Motivated by the conjecture \\citep{CP00} that IRS~13E is an X-ray binary, we present here an analysis of {\\it Chandra} Advanced CCD Imaging Spectrometer (ACIS) GC observations (obs ID=242) of this source. Details of the observations can be found in \\citet{Baganoffetal00,Baganoffetal01b,Baganoffetal01}. The X-ray source $\\sim4\\arcsec$ west-southwest of Sgr A* is the only source in the central parsec that has been associated with a previously known stellar object, IRS~13E. Since IRS~13E appears as an X-ray source while the other early-type stars in the central parsec do not, IRS~13E must harbor a distinctive object. Based on its lack of significant variability at all wavelengths and its strong X-ray luminosity with characteristic $kT \\simeq 1.0$ keV, we argue that IRS~13E is most likely an early-type wide binary system with the primary only recently having exited the luminous blue variable (LBV) phase of evolution. ", + "conclusions": "\\begin{figure} \\resizebox{\\hsize}{!}{\\rotatebox{-90}{\\includegraphics{fig1.ps}}} \\caption[]{A plot of counts per second per keV versus energy for the data (crosses) and best-fitting model (solid line). } \\label{fig:1} \\end{figure} Fig.~\\ref{fig:1} shows the ACIS data along with the best-fitting model, which yielded $\\chi^2_\\nu \\simeq 1.5$, formally a somewhat poor fit. The best-fitting parameters are listed in Table~\\ref{tab:fit}. Inspection of Fig.~\\ref{fig:1} shows that our model underestimates the evident line emission at $\\simeq 3$ keV and slightly overestimates the emission at $\\simeq 3.75$ keV. Some of this may be due to poor statistics but the strong emission near $3$ keV may also reflect a complex metal abundance. \\begin{table} \\caption{Computed Model for IRS~13E and Arches A2} \\label{tab:fit} \\begin{tabular}{lll} \\hline & IRS~13E & Arches A2$^a$ \\\\ \\hline & \\\\[-7pt] $Z$ &1.0 & 2.6$^b$ \\\\ $L_\\mathrm{x}~(\\mathrm{L}_{\\sun})$ &7 &20 \\\\ $kT$ (keV) &$1.0\\pm0.4$ &$1.0\\pm0.5$ \\\\ $N_\\mathrm{H} (10^{22}$ cm$^{-2})$ &$15\\pm5$ &$12\\pm2$ \\\\[3pt] \\hline \\end{tabular} a) from \\citet{Yusefzadehetal01b} \\\\ b) averaged Z over Si, S, Fe, Ar, and Ca \\end{table} Our fit suggests $N_\\mathrm{H}~\\simeq~15\\times~10^{22}$~cm$^{-2}$. This is somewhat larger than the $5$--$10\\times~10^{22}$~cm$^{-2}$ found for Sgr A* \\citep{Baganoffetal00,Baganoffetal01b}, implying that there is additional substantial absorption close to the IRS~13E X-ray source. If so, and if the size of the X-ray source is comparable to the $\\simeq 0.1\\arcsec$ diameter region seen in mm observations, then the characteristic density of the gas surrounding IRS~13E is consistent with a colliding wind binary (CWB) system \\citep{SBP92} but somewhat more dense than LBV ejecta \\citep{Sewardetal01}. However, some of our estimated column to the GC may be due to dust \\citep{Baganoffetal01b} since, unlike most of the massive stars in the central parsec, the IRS~13 complex is enshrouded by warm dust \\citep{M99}. Similar large amounts of dust are seen around ``cocooned'' stars located in the Quintuplet Cluster that is 50 pc from Sgr A* in projection \\citep{FMM99}; although these stars are probably WCd type stars, this is not yet certain \\citep{MSBFN01}. The fitted characteristic temperature, $kT \\simeq 1.0$ keV, is consistent with the shocked winds of a CWB system such as a WR+O. In contrast, an accretion source with a compact companion such as in a massive X-ray binary (MXRB) system, would typically have a harder spectrum \\citep[{\\sl e.g.},][]{Schlegaletal93}. Also, a solitary massive O-star typically has a characteristic temperature of only $\\sim~0.5$~keV \\citep{CHS89}. In the case of a CWB, the characteristic temperature represents a global average of the hot, shocked colliding winds, while X-ray emission from a MXRB probes the accretion disk. On the other hand, the X-ray emission from LBVs is not well known; only $\\eta$ Car, a suspected binary whose primary is an LBV, has a well-determined spectrum. The emission from $\\eta$ Car is a combination of hard ($kT \\sim 5$ keV) compact emission from the star and softer ($kT \\sim 0.5$ keV) extended emission from ejecta \\citep{Sewardetal01}. However, as a binary, the X-ray emission from $\\eta$ Car may not be typical of solitary LBV stars. The absorption-corrected X-ray luminosity of IRS~13E in the 0.2-10 keV band is found to be $L_\\mathrm{x} \\simeq 7$~L$_{\\sun}$. Although larger luminosities are found for the brightest early-type binary systems \\citep[e.g. WR 140 and, assuming binarity, $\\eta$ Car;][] {Huchtetal94,CISP01}, $L_\\mathrm{x} \\simeq 7$~L$_{\\sun}$ is still brighter than the typical CWB system. However, this is not wholly unexpected given the large mass-loss rate and wind velocity of IRS~13E2 \\citep{Najarroetal97}. In contrast, such an X-ray luminosity is a bit low for a MXRB with an accreting compact source unless the system has an unusually long period. Additionally, the X-ray luminosity of a MXRB generally varies on short time-scales. Therefore, due to its low luminosity, lack of variability, and relatively soft spectrum, it is somewhat unlikely that IRS~13E contains a short period MXRB. The lack of variability also makes it unlikely that the X-ray emission from IRS~13E is due to either a flaring proto-star or young stellar object (YSO). The estimated intrinsic X-ray luminosity is also considerably higher than any known YSOs \\citep[{\\sl e.g.},][]{Garmireetal00}. However, \\citet{Clenetetal01} show that the K-L colour of the IRS~13 complex as a whole is consistent with the presence of YSOs. Single massive stars seem to obey the rough relation $\\mathrm{log_{10}} \\left(L^\\mathrm{ISM}_\\mathrm{x}/L_\\mathrm{bol}\\right) = -7\\pm1$ \\citep{Pallavicinietal81}, although more recent work \\citep{CG91,Moffatetal01} suggests the scatter may be more substantial and the ratio may be half a dex larger. $L^\\mathrm{ISM}_\\mathrm{x}$ is the X-ray luminosity corrected for extinction due to the ISM but not for intrinsic extinction \\citep{WCDS98}. If we assume that $N^\\mathrm{ISM}_\\mathrm{H} = 10^{23}$ cm$^{-2}$, then our model for IRS~13E results in $\\mathrm{log_{10}} \\left(L^\\mathrm{ISM}_\\mathrm{x}/L_\\mathrm{bol}\\right) = -7$, no greater than solitary X-ray sources. CWB systems tend to have an enhanced X-ray to bolometric luminosity ratio by a factor of a few compared to solitary stars, but there is considerable variation. Possible contamination and large error bars in determining $\\mathrm{log_{10}} \\left(L^\\mathrm{ISM}_\\mathrm{x}/L_\\mathrm{bol}\\right)$ for IRS~13E2 make it difficult to draw any conclusions concerning binarity using this ratio. For example, it may be that very little of the extinction is intrinsic to the IRS~13E system (see below); this would result in a substantially larger $L^\\mathrm{ISM}_\\mathrm{x}$. We must caution that given the low signal-to-noise of the data, the best-fitting parameters are not very well constrained. For example, forcing $kT = 5$ keV and varying only $N_\\mathrm{H}$ and $L_\\mathrm{x}$ results in a fit with $\\chi_\\nu^2 = 1.9$. However, based on observations of other objects in the GC it is probable that $N^\\mathrm{ISM}_\\mathrm{H} \\simgt 10^{23}$ cm$^{-2}$. Given this additional constraint, we can say that $L_\\mathrm{x} \\simgt 1 \\mathrm{L}_{\\sun}$ and $kT \\simlt 1.8$ keV. Also shown in Table~\\ref{tab:fit} are the results for the bright soft component of source A2 of the Arches cluster located $\\simeq 25$ pc from Sgr A* \\citep{Yusefzadehetal01,Yusefzadehetal01b}. The fit to this source implies a higher than solar metallicity for a range of metals. Although we assume solar metallicity for IRS~13E, the peaks near $3$ and $4.25$ keV in Fig.~\\ref{fig:1} are possible indicators of high S and Ar content. Enhanced abundances for these elements lead to a slightly better fit ($\\chi_\\nu^2 = 1.3$) but due to the low number of counts in our spectrum, we do not attribute much significance to this. Our column density, characteristic temperature, and X-ray luminosity are close to that for A2, which is also coincident with a known IR stellar source, suggesting the objects are of similar nature. In short, even given the large uncertainties, the X-ray luminosity and characteristic temperature of IRS~13E do not favour a single object (O-star, WR star, LBV, or YSO). Also, the lack of variability is inconsistent with a short-period MXRB. On the other hand, all of the characteristics of IRS~13E are fully consistent with a long-period CWB system." + }, + "0112/astro-ph0112053_arXiv.txt": { + "abstract": "We present high-resolution radio and X-ray studies of the composite supernova remnant G11.2$-$0.3. Using archival VLA data, we perform radio spectral tomography to measure for the first time the spectrum of the shell and plerion separately. We compare the radio morphology of each component to that observed in the hard and soft Chandra X-ray images. We measure the X-ray spectra of the shell and the emission in the interior and discuss the hypothesis that soft X-ray emission interior to the shell is the result of the expanding pulsar wind shocking with the supernova ejecta. We also see evidence for spatial variability in the hard X-ray emission near the pulsar, which we discuss in terms of ion mediated relativistic shocks. ", + "introduction": "The supernova remnant G11.2$-$0.3 is a bright, circular X-ray and radio shell. At its center is the X-ray pulsar PSR J1811$-$1925 ($P=65$~ms, spin-down energy $\\dot E = 6.4 \\times 10^{36}$~erg/s, characteristic age $\\tau=24,000$~yrs, Torii et al. 1999) and its associated hard X-ray wind nebula (Vasisht et al. 1996). The characteristic age is much greater than the apparent age of the SNR ($\\sim 2000$~yrs), the age implied by its highly centralized position within the remnant, and its likely association with the historical event of 386~A.D. (Kaspi et al. 2001). This discrepancy suggests the pulsar's current spin period is very near its initial value, and that $\\dot E$ has remained nearly constant since the supernova explosion. ", + "conclusions": "The separation between the PWN and the inner edge of the shell in both radio and X-rays suggests that the remnant may not have reached the Sedov phase. This would imply the supernova ejecta within the shell, which interacts with the PWN, is still in the free expansion phase. The shell itself, however, has probably swept up enough mass to be between the free expansion and Sedov solutions. The expansion of the PWN into the ejecta is supersonic, and should set up a forward shock. This heats the ejecta, possibly resulting in the regions of thermal emission seen interior to the shell. The apparent symmetry around the pulsar of this soft ``PWN\" emission (region 4), along with the relatively greater abundances of heavy elements compared with the bright shell, support this picture. The radio PWN on the one side is sandwiched between the bright, narrow, jet-like hard X-ray feature and the soft ``PWN\" region, which also suggests a connection. Simple spherical models of PWN expansion (eg. Reynolds \\& Chevalier 1984) into the freely expanding interior of an SNR suggest that the ratio of the PWN radius to the SNR radius should be $\\sim 0.2$, assuming that the shell is also freely expanding into the ISM. The larger observed ratio of $\\sim 0.3$ is expected if the shell has slowed due to accumulated ISM mass but the reverse shock has not yet encountered the PWN and compressed it. The bright spots in the hard X-ray nebula may be the equivalent of the Crab's ``wisps\". The motion of wisps may be interpreted in the framework of ion mediated relativistic shocks (Gallant \\& Arons 1994). We can estimate $\\sigma$ (Kennel \\& Coroniti 1984), the ratio of Poynting to particle momentum flux in the wind from the expansion velocity of the PWN (assuming $d\\sim 5$~kpc and an age of 1600~yrs) as $\\sigma \\sim V/c \\sim 0.002$. We estimate the magnetic field in the shocked flow $B \\simeq 3 (\\sigma \\dot{E} /( 4 \\pi r_s^2 c))^{1/2} = 3 \\mu G$ if we assume the position of the wind termination shock $r_s$ is approximately that of spot 1. If we identify the distance between the two bright spots, with the ion gyration radius $r_{L,i}$, we find the ions have passed through $\\sim 3 \\sigma^{1/2} r_{L,i} / r_s \\sim 0.1$ of the total potential drop through the open magnetosphere, similar to the Crab's value of 0.3. We conclude that both the geometry and dynamics (nearly relativistic motion) of the bright spots observed near PSR J1811$-$1925 are consistent with the ion mediated relativistic shock model of Gallant and Arons." + }, + "0112/astro-ph0112265_arXiv.txt": { + "abstract": " ", + "introduction": "Takamizawa discovered an unusual variable star TmzV17 in the course of his photographic patrol of variable stars (Takamizawa 1997). The variable (= GSC 5684.522, J2000.0 coordinates: 18\\h 12\\m 22\\fs 13 -11\\deg 40\\arcm 07\\farcs 1) was later found to be identical with an emission line object AS 289 (Merrill and Burwell 1950). AS 289 has been known as a symbiotic star (Sanduleak and Stephenson 1973; Allen 1978) showing moderately strong He II line and TiO absorption. The object was also studied spectrophotometrically by Blair et al. (1983). Despite its relatively abundant history of spectroscopy, neither outbursts nor variablity has been reported. A typical example of photometry of AS 289 is {\\it V}=13.62 and {\\it B-V}=2.10 (Munari et al. 1992b). The He II to H$\\beta$ line ratio suggests a hot ($\\sim 1.5\\times 10^5$ K) ionizing source (Kenyon 1986). ", + "conclusions": "The present observation is the first that revealed the significant variation of the symbiotic star AS 289. The variation is globally characterized by episodic outbursts observed in 1984 and 1995. This interval of outbursts is typical for those of classical symbiotic variables (the limited observations between these two outbursts may have missed some other brightenings, though). However, the early 1996 fading episode is unique, and remarkably similar to eclipses seen in FG Ser = AS 296 (Munari et al. 1992) and V1413 Aql = AS 338 (Wakuda 1988; Munari 1992). Figure 2 shows the detailed light curve of the current outburst. We suspect the eclipse is responsible for the fading episode of AS 289 during outburst. The profile of light curve suggests the eclipse may be grazing total, while the brightness at minimum is considerably brighter than the pre-outburst level, implying that a certain fraction of the outbursting source remain uneclipsed. It is not still certain regarding the periodicity, but considering the duration of the fading ($\\sim$ 100 d) comparable to that of FG Ser (Munari et al. 1992), the orbital period may be an order of $\\sim$ 600-700 d. This period may imply that the relatively faint level in late 1997 may be somehow related to eclipse, or that the next eclipse was unfortunately missed during the solar conjunction period. \\begin{figure} \\begin{center} \\includegraphics[angle=0,height=6.5cm]{fig2.eps} \\caption{Light curve of AS 289 in 1995--1998} \\label{fig-2} \\end{center} \\end{figure}" + }, + "0112/astro-ph0112115_arXiv.txt": { + "abstract": "Recent analysis of the Doppler shift oscillations of the light from extra-solar planetary systems indicate that some of these systems have more than one large planet. In this case it has been shown that the masses of these planets can be determined without the familiar ambiguity due to the unknown inclination angle of the plane of the orbit of the central star provided, however, that its mass is known. A method is presented here which determines also a lower limit to the mass of this star from these observations. As an illustration, our method is applied to the Keck and Lick data for GJ876. ", + "introduction": " ", + "conclusions": "" + }, + "0112/astro-ph0112337_arXiv.txt": { + "abstract": "Kaastra et al. (1999) have used the BeppoSAX LECS instrument to search for excess EUV emission in Abell\\ 2199. They claim that the results obtained confirm an independent report of an excess EUV emission in this cluster (Lieu et al. 1999). Using an inflight derived procedure that is better suited to the analysis of extended sources and which avoids uncertainties related to ground-based calibrations for the overall detector sensitivity profile, we find no excess EUV emission in Abell\\ 2199. We also used these procedures to search for an EUV excess in Abell\\ 1795, but no excess was found. ", + "introduction": "\\label{intro} The discovery of Extreme Ultraviolet (EUV) emission in clusters of galaxies with the Extreme Ultraviolet Explorer (EUVE) has provoked considerable controversy. While there is no doubt about the detection of the EUV emission in excess of that produced by the well-studied X-ray emitting cluster gas in Virgo (Bergh\\\"ofer, Bowyer, \\& Korpela 2000a) and in Coma (Bowyer, Bergh\\\"ofer, \\& Korpela 1999), many clusters do not exhibit an EUV excess at least at current sensitivity levels. In a series of publications (Bowyer, Bergh\\\"ofer, \\& Korpela, 1999; Bergh\\\"ofer, Bowyer, \\& Korpela 2000a,b; Bowyer, Korpela, \\& Bergh\\\"ofer 2001) we have demonstrated that the EUVE results are strongly affected by the variation of the telescope sensitivity over the field of view and upon the details of the subtraction of the EUV emission from the X-ray contribution. Kaastra et al. (1999) have analyzed BeppoSAX data obtained with the Low-Energy Concentrator Spectrometer (LECS) to search for EUV emission in the Abell\\ 2199 cluster of galaxies. Unfortunately, the telescope sensitivity profile used in this work is likely to be incorrect. Quoting from their work, Kaastra et al. (1999) state, \"The vignetting correction for the LECS was derived from the SAXDAS/LEMAT ray-trace code, assuming azimuthal symmetry around the appropriate center. The correction for the support grid was also derived from that package.\" We note that ground based simulations of the large scale sensitivity of EUV and X-ray instrumentation are extraordinarily difficult to construct and are notorious for being an inappropriate representation of the true sensitivity functions. In fact, a comparison of ground based calibration measurements with ray-tracing simulations for BeppoSAX clearly demonstrated a discrepancy by a factor of $\\sim$1.5 at low energies (Parmar et al. 1997). In order to test the claim of an EUV excess in Abell\\ 2199 by Kaastra et al. (1999) we have reanalysed BeppoSAX LECS observations of this cluster of galaxies. We also searched for excess EUV emission in Abell\\ 1795 using archival BeppoSAX data on this cluster. In Section\\ \\ref{data} we describe the data and reduction applied to obtain the clusters' EUV emission. Section\\ \\ref{results} provides the results of our investigations, and a summary is presented in Section\\ \\ref{discussion}. ", + "conclusions": "\\label{discussion} A search for excess EUV emission has been carried out in a substantial number of clusters observed with EUVE. Bowyer, Bergh\\\"ofer \\& Korpela (1999) have shown that all but two of the reported detections were the product of the use of an incorrect detector sensitivity function. However, this conclusion has been questioned because Kaastra et al. (1999) claimed to have found an excess EUV emission in Abell 2199 using BeppoSAX LECS observations. These findings appeared to support the (incorrect) finding of an excess in this cluster using EUVE data. However, this result was based upon ground-based estimates of the LECS detector sensitivity function. Using a procedure better suited to the analysis of extended sources that avoids the known uncertainties in the telescope sensitivity function, we show that there is no excess in Abell\\ 2199. We also searched for excess EUV emission in Abell\\ 1795 using archival BeppoSAX observations. No EUV excess was found. The results obtained here for Abell\\ 2199 and Abell\\ 1795 are fully consistent with the results obtained on these clusters using EUVE data (Bowyer, Bergh\\\"ofer, \\& Korpela 1999). The only clusters of galaxies that have been found to exhibit an excess EUV emission are the Virgo and Coma clusters (Bowyer, Korpela, \\& Bergh\\\"ofer 2001). It is possible that these are the only clusters that exhibit this effect, but it may be that both EUVE and BeppoSAX LECS are insufficiently sensitive to detect an EUV excess in other clusters of galaxies. Since the only clusters with a confirmed EUV excess are Virgo and Coma, it is useful to reconsider candidates for the underlying source of the EUV excess. The original proposal was that this emission is thermal emission from a ``warm'' (10$^6$ K) gas (Lieu et al. 1996a,b; Bowyer et al. 1996). Claims of ``proof'' of this proposition have been advanced by Mittaz, Lieu, \\& Lockman (1998), Lieu, Bonamente, \\& Mittaz (1999a), Lieu et al. (1999b), Lieu, Bonamente, \\& Mittaz (2000), and Bonamente, Lieu, \\& Mittaz (2001a,b). Lieu et al. (2000) misinterpreted small scale detector structures in the EUVE data on Abell 2199 as cluster EUV emission absorbed by clumps of neutral hydrogen in the cluster. A difficulty with this interpretation which is independent of the data analysis problem is that the hydrogen required is \"... $\\sim$43 times more massive than the hot ICM in this region ... (and) ... implies 3 times more missing baryons than expected''(op. cite). The maintenance of a warm intracluster gas is quite difficult to understand since gas at this temperature is at the peak of its cooling curve and would cool in less than 0.5~Gyr, and on these grounds alone it was generally believed that a thermal source was untenable. Observational evidence relevant to this issue was obtained with the Hopkins Ultraviolet telescope (Dixon et al. 1996), and FUSE (Dixon et al. 2001a,b). No Far UV line emission from gas at 10$^6$ K was detected. More recently, observations of a large number of clusters with XMM have been carried out. Kaastra et al. (2001) found no gas at T $<$ 1 keV in Sersic 159-03, Peterson et al. (2001) found no gas at T $<$ 2.7 keV in Abell 1835, and Tamura et al. (2001) found no gas at T $<$ 4 keV in Abell 1795. All other clusters observed with XMM showed no evidence of a cooler EUV emitting gas (Steve Kahn, private communication). The sum of this evidence seems overwhelming: a thermal mechanism for the EUV excess can be ruled out. Since the underlying source mechanism is not thermal, it must be the product of some non-thermal process. Inverse Compton scattering of cosmic rays with the 2.7 K background was suggested early-on as a possible source mechanism (Hwang 1997, En{\\ss}lin \\& Biermann 1998). Sarazin \\& Lieu (1998) suggested a model in which a population of cosmic rays produced several Gyr ago would have degraded over time and would now be unobservable as radio synchrotron emission even at very low frequencies. This population would produce an EUV flux by inverse Compton scattering. Sarazin \\& Lieu derived the ratio between the azimuthally averaged total EUV emission and the azimuthally averaged soft X-ray flux predicted by their model; this ratio increases with increasing distance from the center of the cluster. Bergh\\\"ofer et al. (2000) derived this ratio for the Virgo cluster as a test of the Sarazin \\& Lieu model. They found this ratio was flat with increasing distance from the center of the cluster in contradiction to the prediction of the model. Bowyer et al. (in progress) derived this ratio for the Coma cluster using data on the cluster that had been analyzed correctly. They found this ratio was flat with increasing distance from the center again contradicting the predictions of the Sarazin \\& Lieu model. Despite the failure of the Sarazin \\& Lieu model, the inverse Compton mechanism remains as the only candidate for the source mechanism for the EUV excess. However, a new difficulty for this hypothesis has recently appeared. Virtually all models invoking the inverse Compton mechanism require the intracluster magnetic field to be $<<$1$\\mu$G. However, recent results show that cluster magnetic fields are quite large. Clarke, Kronberg, \\& B\\\"ohringer (2001) studied 16 clusters with very high spatial resolution and have shown that all of these clusters have B fields of 4 to 7$\\mu$G. Unless this result is somehow incorrect, the vast majority of models proposed for the production of the EUV excess are incorrect. The only exceptions to the low field models (En{\\ss}lin, Lieu, \\& Biermann 1999; Atoyan \\& V\\\"olk 2000) are unlikely to be appropriate (Ming, Hwang, \\& Bowyer 2001). It is not clear whether this is a fundamental obstacle for the inverse Compton scattering hypothesis or if it is simply a failure of existing models. Irrespective of how widespread the occurrence of EUV excess in clusters of galaxies may be, the underlying source mechanism for this emission remains a mystery." + }, + "0112/astro-ph0112101_arXiv.txt": { + "abstract": "Optical spectroscopic observations of the companion star (type G8IV to K4III) in the microquasar system XTE J1550-564 reveal a radial velocity curve with a best fitting spectroscopic period of $P_{\\rm sp}=1.552\\pm 0.010$ days and a semiamplitude of $K_2=349\\pm 12$ km s$^{-1}$. The optical mass function is $f(M)=6.86\\pm 0.71\\,M_{\\odot}$ ($1\\sigma$). We tentatively measure the rotational velocity of the companion star to be $V_{\\rm rot}\\sin i=90\\pm 10$ km s$^{-1}$, which when taken at face value implies a mass ratio of $Q\\equiv M_1/M_2=6.6^{+2.5}_{-1.6}$ ($1\\sigma$), using the above value of $K_2$. We derive constraints on the binary parameters from simultaneous modelling of the ellipsoidal light and radial velocity curves. We find $1\\sigma$ ranges for the photometric period ($1.5430\\,{\\rm d}\\le P_{\\rm ph}\\le 1.5440\\,{\\rm d}$), $K-$velocity ($350.2\\le K_2\\le 368.6$ km s$^{-1}$), inclination ($67.0^{\\circ}\\le i\\le 77.4^{\\circ}$), mass ratio ($Q\\ge 12.0$), and orbital separation ($11.55\\,R_{\\odot}\\le a\\le 12.50\\,R_{\\odot}$). Given these geometrical constraints we find the most likely value of the mass of the compact object is $9.41\\,M_{\\odot}$ with a $1\\sigma$ range of $8.36\\,M_{\\odot}\\le M_1\\le 10.76\\,M_{\\odot}$. If we apply our tentative value of $V_{\\rm rot}\\sin i=90\\pm 10$ km s$^{-1}$ as an additional constraint in the ellipsoidal modelling, we find $1\\sigma$ ranges of $1.5432\\,{\\rm d}\\le P_{\\rm ph}\\le 1.5441\\,{\\rm d}$ for the photometric period, $352.2\\le K_2\\le 370.1$ km s$^{-1}$ for the $K-$velocity, $70.8^{\\circ}\\le i\\le 75.4^{\\circ}$ for the inclination, $6.7\\le Q\\le 11.0$ for the mass ratio, and $12.35\\,R_{\\odot}\\le a\\le 13.22\\,R_{\\odot}$ for the orbital separation. These geometrical constraints imply the most likely value of the mass of the compact object of $10.56\\,M_{\\odot}$ with a $1\\sigma$ range of $9.68\\,M_{\\odot}\\le M_1\\le 11.58\\,M_{\\odot}$. In either case the mass of the compact object is well above the maximum mass of a stable neutron star, and we therefore conclude XTE J1550-564 contains a black hole. ", + "introduction": "X-ray novae provide the strongest evidence for the existence of stellar mass black holes. These objects are interacting binaries containing a compact primary (a neutron star or a black hole) and what is usually a late-type secondary. X-ray novae spend most of their time in a low X-ray luminosity ``quiescent'' state, with an X-ray luminosity ($L_x$) roughly similar to the optical luminosity ($L_{\\rm opt}$), which includes a substantial contribution from the secondary star ($L_{\\rm star}$). The quiescent state is occasionally interrupted by ``outbursts'' where typically $L_x\\gg L_{\\rm opt}\\gg L_{\\rm star}$. In quiescence, the observed radial velocity and light curves of the secondary star lead to dynamical mass estimates for the compact primary. If its mass exceeds the maximum stable mass of a neutron star ($\\approx 3\\,M_{\\odot}$), the compact object is presumed to be a black hole \\citep{chi76,kal96}. In thirteen cases, the mass of the primary of an X-ray nova has been shown to exceed $3\\,M_{\\odot}$, confirming the presence of black holes in these systems \\citep{bai95,cas92,cas95,fil95a,fil99,fil01,gre01, mcc86,mcc01,oro98,oro01,rem92, rem96}. These X-ray novae open up the possibility of studying the strong-field regime of general relativity. For example, quiescent X-ray spectra and ADAF (advection-dominated accretion flow) models provide evidence for black hole event horizons \\citep{nar96,gar01}. Also, the study of high-frequency QPOs (quasi-periodic oscillations) may lead to the first secure measurement of black hole spin (e.g.\\ Remillard et al.\\ 1999; Strohmayer 2001). XTE J1550-564 was discovered on 1998 September 7 by the All-Sky Monitor (ASM) aboard the {\\em Rossi X-ray Timing Explorer (RXTE)} \\citep{smi98}. The optical counterpart (designated V381 Normae) and the radio counterpart were discovered shortly thereafter \\citep{oro98a,cam98}. This source was quickly identified as a promising black hole candidate based on its rapid X-ray variability, hard energy spectrum, and the absence of pulsations or X-ray bursts \\citep{cui99,sob99}. Relativistic plasma ejections at probable superluminal velocities were observed at radio wavelengths shortly after the strong X-ray flare in 1998 September \\citep{han01b}, indicating that XTE J1550-564 is another microquasar. The galactic microquasars are excellent laboratories for the study of relativistic jets since they evolve orders of magnitude more quickly than do the jets in quasars \\citep{mir99}. XTE J1550-564 is also of special interest owing to its complex X-ray variability (e.g.\\ Homan et al.\\ 2001; Remillard et al.\\ 2002). In this paper we report the results of our recent optical observations of XTE J1550-564. The observations and basic data reductions are summarized in \\S\\ref{obssec}. In \\S\\ref{ansec} we establish the orbital parameters of the system, derive some properties of the secondary star, derive geometrical parameters for the binary, and discuss limits on the mass of the compact object. The implications of our results are discussed in \\S\\ref{dissec} and summarized in \\S\\ref{sumsec}. ", + "conclusions": "\\label{dissec} \\citet{bai98} analyzed the distribution of black hole masses in seven systems and found a high probability that six of the seven systems (GRO J0422+32, A0620-00, GS 1124-683, GRO J1655-40, H1705-250, and GS 2000+25) have masses which are consistent with $\\approx 7\\,M_{\\odot}$. The seventh system, V404 Cyg, has a mass which is significantly larger (about $12\\,M_{\\odot}$; Shahbaz et al.\\ 1994). The mass of the black hole in XTE J1550-564 is well above $7\\,M_{\\odot}$. With $M_1\\ge 8.1\\,M_{\\odot}$ at $3\\sigma$ confidence (Table \\ref{parm}), its mass may in fact be similar to the mass of the black hole in V404 Cyg. Since the work of \\citet{bai98}, the mass functions for seven additional systems have been measured, and improved parameters for some of the original seven systems have been measured. Thus the issue of the observed black hole mass distribution should be revisited to see if the clustering of black hole masses near $7\\,M_{\\odot}$ is still significant. Given the new data, we can begin to make meaningful comparisons with formation theory. For example, the detailed formation models of \\citet{fry01} predict a mass distribution which is continuous and which extends over a broad range (in particular, they predict no peak at $\\approx 7\\,M_{\\odot}$). The determination of the black hole mass for XTE J1550-564 is especially important for the interpretation of the high-frequency X-ray quasi-periodic oscillations (QPOs) observed for this system (Remillard et al. 2002; Homan et al. 2001). Models for several types of oscillations predicted in general relativity are under investigation as possible causes of these QPOs (e.g. Remillard 2001, and references therein); all of these depend on both the mass and spin of the black hole, and possibly also on conditions in the inner accretion disk. Despite considerable uncertainties in the models, we can offer a few comments on the implications of our mass determination for XTE J1550-564. At the nominal mass of $10.7\\,M_{\\odot}$, oscillations at the frequency of the last stable orbit for a Schwarzschild black hole (spin parameter, $a_* = 0$) would be seen at a frequency, $\\nu = 2199 / (M_{\\rm BH} / M_{\\odot}) = 208.2 $ Hz \\citep{sha83}. Since XTE J1550-564 has exhibited QPOs with frequencies up to 284 Hz \\citep{hom01}, it appears that plausible mechanisms require $a_* > 0$. (The oscillation frequency would be 284 Hz (for $a_*=0$) if $M_{\\rm BH}=7.74\\,M_{\\odot}$, which is below the $3\\sigma$ lower limit on the black hole mass.) This argument was used by \\citet{str01} to suggest that the 450 Hz QPO in GRO J1655-40 implies appreciable spin for the black hole in that system. We further note that the ratio of the black hole masses for XTE J1550-564 ($M_1=10.6\\,M_{\\odot}$, Table \\ref{parm}) and GRO J1655-40 ($M_1=6.3\\,M_{\\odot}$, Greene, Bailyn, \\& Orosz 2001) is $\\sim 1.7$, which is nearly inversely proportional to their maximum QPO frequencies: $(284 {\\rm Hz}/{\\rm 450 Hz})^{-1}\\sim 1.6$. We may then speculate that these results are consistent with a common QPO origin, with QPO frequencies that vary as $M^{-1}$, which might be expected if the black holes have similar values of the spin parameter. Both XTE J1550-564 and GRO J1655-40 have had several closely-spaced outburst events relatively shortly after their initial discoveries. Interestingly enough, optical precursors to X-ray outbursts (i.e.\\ X-ray delays) have been observed for both sources. According to \\citet{ham97}, these X-ray delays can best be understood using a two component accretion model consisting of a standard thin disk in the outer regions and an ADAF region in the inner region. At the onset of an outburst, the heating front reaches the ADAF region relatively quickly, but then must propagate more slowly on a viscous timescale, thereby greatly delaying the production of X-rays. \\citet{oro97a} observed an optical precursor to the 1996 April outburst of GRO J1655-40. In that case, the delay between the $V$-band rise and the detection by the ASM on {\\em RXTE} was $5.6\\pm 0.8$ days. In the case of the 2000 outburst of XTE J1550-564, \\citet{jai01c} reported a delay of $8.8\\pm 0.6$ days between the rise in $V$ and the detection by the {\\em RXTE} ASM. The ratio of these delays is $8.8/5.6\\sim 1.6$, which is nearly the same as the ratio of the masses and is the same as the inverse ratio of their maximum QPO frequencies ($\\sim 1.7$ and $\\sim 1.6$ respectively). We may again speculate that these results imply that the size of the ADAF cavity scales with the mass of the black hole. For the case of a companion star mass of $1.3\\,M_{\\odot}$ (Table \\ref{parm}), the companion star in XTE J1550-564 appears to have mass that is larger than the typical mass for X-ray novae with a cool (K or M) companion. For comparison, the K0 secondary star in V404 Cyg has a mass of about $0.7\\,M_{\\odot}$ \\citep{sha94}. Since the secondary star in XTE J1550-564 is about three times larger than a normal K3/4 dwarf, its evolutionary state may be different than the dwarf secondaries in the short period X-ray novae such as A0620-00 or GS 1124-683. If the star is still in the core hydrogen burning phase, then its nuclear evolution timescale would be $\\approx 1\\times 10^9$ years, leading to an average mass transfer rate of $10^{-9}\\,M_{\\odot}$ yr$^{-1}$. If the star has finished its core hydrogen burning, then presumably its radius is expanding on a thermal timescale of $\\approx 10^{7}$ years, leading to an average mass transfer rate of $10^{-7}\\,M_{\\odot}$ yr$^{-1}$. We can make a rough estimate of the mass transfer rate in XTE J1550-564 as follows. The {\\em RXTE} satellite performed numerous observations of XTE J1550-564 during its 1998 and 1999 outbursts. \\citet{sob00} give the unabsorbed 2-20 keV flux and the 20-100 keV flux for the $\\approx 200$ PCA and HEXTE observations from this period, respectively (their Fig.\\ 6). Integrating under the curves, we find a total 2-100 keV fluence of 1.06 erg cm$^{-2}$ for the 1998-1999 outburst. Integrating the ASM light curve (Fig.\\ 1) we find that the outburst in 2000 was about 20 times fainter in the 2-12 keV band than the 1998 and 1999 events. Assuming similar spectral characteristics for the two events, we conclude that total X-ray fluence in the 2-100 keV band was 1.1 erg cm$^{-2}$ in the interval between 1998 and 2001. To get the average mass transfer rate, we need to know the recurrence time between major outbursts and the distance. According to Chen, Shrader, \\& Livio (1997), the sky coverage provided by scanning X-ray instruments has been quite good since about 1988, so it seems likely that the recurrence time of XTE J1550-564 is at least ten years. The distance is not well known, but it is probably about 6 kpc to within a factor of two (Table \\ref{tabdist}). In Table \\ref{tabenergy} we give total isotropic energies and average mass transfer rates for assumed distances of 3, 6, and 10 kpc and recurrence times of 10 and 50 years (we assume an accretion efficiency of 10\\% and that the mass transfer rate from the secondary star is the same as the mass transfer rate onto the black hole). The mass transfer rates are between about $1\\times 10^{-10}$ and $7\\times 10^{-9}\\,M_{\\odot}$ yr$^{-1}$, which suggests the star is evolving on a relatively long nuclear timescale. GRO J1655-40 might be a similar system in terms of its evolutionary state. The average mass transfer rate in GRO J1655-40 has been estimated to be about $1\\times 10^{-10}\\,M_{\\odot}$ yr$^{-1}$ (van Paradijs 1996). However, the secondary star is located in the Hertzsprung gap in the Hertzsprung-Russell diagram (e.g.\\ Orosz \\& Bailyn 1997), so one would expect a mass transfer rate of about $1\\times 10^{-7}\\,M_{\\odot}$ yr$^{-1}$ (e.g.\\ Kolb et al.\\ 1997). Regos, Tout, \\& Wickramasinghe (1998) have suggested that the secondary star in GRO J1655-40 is still in the core hydrogen burning phase, and they have constructed evolutionary models with mass transfer rates much closer to the observed value. It remains to be seen if a similar model can be constructed for XTE J1550-564. Although the presence of a black hole in XTE J1550-564 is now firmly established, follow-up observations would be desirable to confirm and improve upon our results. The statistical error in the mass function can be easily reduced by adding more radial velocities. Spectroscopic observations with higher resolution should be obtained so that the rotational velocity of the companion star can be measured more precisely. Our inclination limits can be improved by obtaining multicolor light curves with much better phase coverage." + }, + "0112/astro-ph0112047_arXiv.txt": { + "abstract": "A neutron star with mass close to the lower limit might be a reasonable model for some anomalous pulsars. Emission is thermal. X-ray luminosity is high. Spatial velocity can be high. Since the radius is predicted to be large, the magnetic field calculated for spin-down is lower than that required by the magnetar model. ", + "introduction": "Almost every pulsar has precisely measured period, $P$, and the rate of change of period with time, $\\dot{P}$. The gross properties of these objects are well illustrated by the $P-\\dot{P}$ diagram which now shows rotational data for $\\sim 1200$ isolated (as far as we know) pulsars. The accepted model, a rotating, strongly magnetized, neutron star, gives a plausible explanation for the behavior and evolution of most of these objects. Approximately 10 objects, having both large $P$ and $\\dot{P}$, form a loose group in the upper right corner. Here the dipole model predicts large values for the magnetic moment, $M$, and the surface field, $B$. Indeed, the energy in the predicted magnetic field exceeds the rotational energy in this part of the diagram. These objects are the Anomalous X-ray Pulsars (AXP) and the the Soft Gamma-ray Repeaters (SGR) which we refer to collectively as anomalous pulsars. Also for these pulsars the observed X-ray luminosity also greatly exceeds the rate of loss of rotational energy. Some anomalous pulsars are associated with supernova remnants, and the short lifetimes indicated by $P/2\\dot{P}$ seem appropriate. It has been proposed (Thompson \\& Duncan 1996) that energy is supplied by decay of the very-strong magnetic field. The interaction of the strong field with the crust of the neutron star is also invoked to explain the $\\gamma$-ray bursts observed from SGR. AXP/SGRs are also called ``magnetars'' in the literature. It is assumed that the slowing torque is largely electromagnetic. However, the grouping of anomalous pulsars in the upper right of the $P-\\dot{P}$ diagram, suggests that something different is happening rather than an extension of the dipole model. We here explore the possibility that the anomalous pulsars are low-mass neutron stars -- objects with mass considerably smaller than the generally-accepted mass of 1.4 M$_\\odot$. ", + "conclusions": "" + }, + "0112/astro-ph0112517_arXiv.txt": { + "abstract": "{A temporal analysis has been performed on a sample of 100 bright gamma-ray bursts (GRBs) with $T_{90}$$<$$2\\,$s from the BATSE Current Catalog. The GRBs were denoised using a median filter and subjected to an automated pulse selection algorithm as an objective way of identifing the effects of neighbouring pulses. The rise times, fall times, FWHM, pulse amplitudes and areas were measured and the frequency distributions are presented here. All are consistent with lognormal distributions. The distribution of time intervals between pulses is not random but consistent with a lognormal distribution. The time intervals between pulses and pulse amplitudes are highly correlated with each other. These results are in excellent agreement with a similar analysis that revealed lognormal distributions for pulse properties and correlated time intervals between pulses in bright GRBs with $T_{90}$$>$$2\\,$s. The two sub-classes of GRBs appear to have the same emission mechanism which is probably caused by internal shocks. They may not have the same progenitors because of the generic nature of the fireball model. ", + "introduction": "It has been recognised that GRBs may occur in two sub-classes based on spectral hardness and duration with $T_{90}$$>$$2\\,$s and $T_{90}$$<$$2\\,$s \\citep{kmf:1993,dbt:1995,paciesas:2001}. The bimodal distribution can be fit by two Gaussian distributions in the logarithmic durations \\citep{mhlm:1994}. There is significant evidence for a third subgroup as part of the long duration GRBs \\citep{mfb:1998,horv:1998} but this has been questioned because of a possible BATSE selection effect \\citep{hhp:2000}. It has been suggested \\citep{cmo:1999} that the small group of GRBs with T$_{90} <$ 0.1 s form an additional category. The short GRBs have a higher value of $\\langle V/V_{\\rm max}\\rangle$ \\citep{kc:1996}, a much smaller value of the spectral lag \\citep{nsb:2000}, a pulse shape that depends on position in the burst \\citep{gupta:2000} and a smaller space density than long GRBs \\citep{schmidt:2001}. A variety of statistical methods have been applied to the temporal properties of GRBs with $T_{90}$$>$$2\\,$s. It is important to compare the temporal profiles of the long and short GRBs to determine the similarities and differences between the two classes in an objective way. A detailed objective analysis has been performed on the temporal profiles of a large sample of 319 bright GRBs with $T_{90}$$>$$2\\,$s \\citep{quillig:2001,hmq:1998}. The properties of the pulses in GRBs and the time intervals between them were found to be consistent with lognormal distributions. These results can be used as templates for comparison with a similar analysis of GRBs with $T_{90}$$<$$2\\,$s. The analysis method is presented in section 2 and the results in section 3. In section 4 the results are discussed and compared with the sample of long GRBs. \\begin{table*}[htbp] \\caption{ The parameters of the pulse properties in GRBs include the median value for the data, the median $\\mu$ and the standard deviation $\\sigma$ expressed as natural logarithms, the width of the lognormal distributions at the $\\pm50\\%$ level in normal space, the KS probability that properties of pulses selected at different levels of $\\tau_{\\rm \\sigma}$ and $\\tau_{\\rm i}$ were drawn from the same distribution as those selected at $\\mathrm{\\tau}_{\\rm \\sigma}$ $\\geq$ 5, $\\mathrm{\\tau}_{\\rm i}$ $\\geq$ 50\\%. All pulses with $\\mathrm{\\tau}_{\\rm \\sigma}$ $\\geq$ 5 were used for the time intervals and not just isolated pulses. The values for the pulse areas and amplitudes were obtained from 55 GRBs that were summed over two BATSE Large Area Detectors. % } \\setlength{\\tabcolsep}{0.3cm} \\begin{tabular}[b]{lccccccc} \\vspace{0.1cm} Property & Median & $\\mu$ & $\\sigma$ & Width ($\\pm50\\%$) & KS ($\\mathrm{\\tau}_{\\rm \\sigma}$ $\\geq$ 3,$\\mathrm{\\tau}_{\\rm i}$ $\\geq$ 20\\%) & KS ($\\mathrm{\\tau}_{\\rm \\sigma}$ $\\geq$ 8, $\\mathrm{\\tau}_{\\rm i}$ $\\geq$ 80\\%) & \\\\ \\hline Rise Time & 0.035 & $-$3.31 & 0.94 & 0.012-0.11 & 0.48 & 0.16 & \\\\ Fall Time & 0.056 & $-$2.89 & 0.98 & 0.017-0.176 & 0.07 & 0.19 & \\\\ FWHM & 0.045 & $-$3.17 & 1.01 & 0.013-0.138 & 0.73 & 0.48 & \\\\ Area & 1.5$\\times 10^{5}$& 11.7 & 1.24 & 28-520($\\times10^{3}$) & 0.44 & 0.15 & \\\\ Pulse Amplitude & 1.02$\\times 10^{4}$ & 9.22 & 0.83 & 3.8-27($\\times10^{3}$) & 0.90 & 0.72 &\\\\ Asymmetry Ratio& 0.65 & $-$0.42 & 0.91 &0.19-2.22 & 0.78 & 0.30 & \\\\ $\\Delta$T & 0.095 & $-$2.24 & 0.87 & 0.038-0.3 & 0.83 & 0.96 & \\\\ \\hline \\end{tabular} \\end{table*} ", + "conclusions": "A sample of bright GRBs with $T_{90}$$<$$2\\,$s have been denoised and analysed by an automatic pulse selection algorithm. The results show that the distribution of the properties of isolated pulses and time intervals between all pulses are compatible with lognormal distributions. The same mechanism seems to be responsible for both long and short GRBs and is attributed to the internal shock model." + }, + "0112/astro-ph0112384_arXiv.txt": { + "abstract": "We consider radiative electromagnetic corrections, at order $\\alpha$, to the process $e^+ e^- \\rightarrow \\nu \\overline{\\nu}$ at finite density and temperature. This process represents one of the main contributions to the cooling of stellar environments in the late stages of star evolution. We find that these corrections affect the energy loss rate by a factor $(-4 \\div 1) \\%$ with respect to the tree level estimate, in the temperature and density ranges where the neutrino pair production via $e^+ e^-$ annihilation is the most efficient cooling mechanism. ", + "introduction": "The late stages of star evolution are strongly influenced by neutrino emission processes. After hydrogen burning, stars with small mass, of the order of the solar mass, evolves towards a white dwarf configuration. On the other hand more massive stars undergo several burning phases till the formation of a central iron core and the eventual triggering of the gravitational instability leading to the Supernova phenomenon. In both cases the cooling rate is largely dominated by neutrino production. For example, from the $^{12}C - ^{24}Mg$ burning phase, almost $100 \\%$ of the energy flux is emitted via neutrino production, which leave the stellar system without any interaction because of their extremely large mean free path. A precise determination of neutrino emission rates is therefore a crucial issue in any careful study of the final branches of star evolutionary tracks. In particular, changing the cooling rates at the very last stages of massive star evolution may sensibly affect the evolutionary time scale and the iron core configuration at the onset of the Supernova explosion, whose triggering mechanism is still lacking a full theoretical understanding \\cite{janka}. The energy loss rate due to neutrino emission, hereafter denoted by $Q$, receives contribution from both weak nuclear reactions and purely leptonic processes. However for the rather large values of density and temperature which characterize the final stages of stellar evolution, the latter are largely dominant. They are mainly due to four possible interaction mechanisms: \\begin{itemize} \\item[i)] pair annihilation $~~~~~~~~~~~~~~e^+ \\, + \\, e^- \\; \\rightarrow \\; \\nu \\, + \\, \\ov{\\nu}$ \\item[ii)] $\\nu$-photoproduction $~~~~~~~~~~~~~ \\gamma \\, + \\, e^{\\pm} \\; \\rightarrow \\; e^{\\pm} \\, + \\, \\nu \\, + \\, \\ov{\\nu}$ \\item[iii)] plasmon decay $~~~~~~~~~~~~~~~~~~~~~~~~~\\gamma^\\ast \\; \\rightarrow \\; \\nu \\, + \\, \\ov{\\nu}$ \\item[iv)] bremsstrahlung on nuclei $~~~~e^{\\pm} \\, + \\, Z \\; \\rightarrow \\; e^{\\pm} \\, + \\, Z \\, + \\, \\nu \\, + \\, \\ov{\\nu}$ \\end{itemize} Actually each of these processes results to be the dominant contribution to $Q$ in different regions in a density-temperature plane. For very large core temperatures, $T \\gsim 10^9 \\,^\\circ K$, and not too high values of density, pair annihilations are most efficient, while $\\nu$ photoproduction gives the leading contribution for $10^8 \\,^\\circ K \\lsim T \\lsim$ $10^9\\,^\\circ K$ and relatively low density, $\\rho \\lsim 10^5$ g cm$^{-3}$. These are the typical ranges for very massive stars in their late evolution. Finally, plasmon decay and bremsstrahlung on nuclei are mostly important for large ($\\rho \\gsim 10^{6}$ g cm$^{-3}$) and extremely large ($\\rho \\gsim 10^{9}$ g cm$^{-3}$) core densities, respectively, and temperatures of the order of $10^8 \\,^\\circ K \\lsim T \\lsim 10^{10} \\,^\\circ K$. Such conditions are typically realized in white dwarfs. Starting from the first calculation of neutrino energy loss rates, based on the V-A theory of weak interactions \\cite{Beaudet67a,Beaudet67b}, a systematic study of processes i)-iv) has been carried out in a long series of papers \\cite{Beaudet67a}-\\cite{Itoh96}. In all these analyses the pair production rate i) has been evaluated at order $G_F^2$, i.e. at the zero order in the electromagnetic coupling constant $\\alpha$ expansion. Despite of the different topology and phase space volume which characterize the remaining processes ii)-iv), it is nevertheless worth noticing that they are instead at least of order $\\alpha G_F^2$. In view of this, it is therefore meaningful to investigate whether including QED radiative corrections to pair annihilation rate i) may lead to a sensible change in the cooling rate $Q$. This is the aim of this letter. In particular we report here the results of a calculation at order $\\alpha$ of two classes of contributions, due to vacuum radiative corrections as well as those arising from the interaction of the $e^+$-$e^-$ pairs with the surrounding electromagnetic plasma. This has been performed in the real time formalism for finite temperature field theory. We actually find that the total corrections to $Q$ range in the interval $(-4 \\div 1) \\%$ of the tree level estimate, in the temperature and density ranges where the pair annihilation process represents the main contribution to the total energy loss rate. We have also re-evaluated both photoproduction and plasmon decay rates, and we do find a good agreement with the most recent estimates reported in the literature \\cite{Raffelt95,Itoh96}. The paper is organized as follows. Section 2 is a short review of the lowest order calculation of $Q$ for the pair annihilation process. The order $\\alpha$ QED corrections are then considered in Section 3, where we report the main results of our calculations of both $vacuum$ and $thermal$ contributions. We finally discuss our results and give our conclusions in Section 4. ", + "conclusions": "In this letter we have reported on the calculation of the energy loss rate of pair annihilation process i) up to order $\\alpha G_F^2$. The Born value for $Q$ has been corrected including both $vacuum$ and $thermal$ radiative corrections, the latter being computed in the real time formalism. One of our main results is presented in Fig. \\ref{paircorr}, where we plot these radiative corrections with respect to the Born approximation estimate, as functions of the plasma density for some values of the temperature. The corrections come out to be of the order of few percent and negative for high temperatures, implying that for these temperatures the energy loss is sensibly decreased. At fixed temperature, $\\Delta Q / Q_B$ goes to a constant value for low density. This can be easily understood, since in this limit the plasma is weakly degenerate, and therefore the energy loss rate depends on temperature only. At large densities the ratio $\\Delta Q/Q_B$ decreases and reach larger negative values. However for such high densities the pair annihilation rates are exceedingly small and thus this process gives only a marginal contribution to the star cooling. \\begin{figure} \\begin{center} \\begin{tabular}{ll} \\epsfxsize=7cm \\epsfysize=10cm \\epsffile{Q8.eps} & \\epsfxsize=7cm \\epsfysize=10cm \\epsffile{Q85.eps} \\\\ \\epsfxsize=7cm \\epsfysize=10cm \\epsffile{Q9.eps} & \\epsfxsize=7cm \\epsfysize=10cm \\epsffile{Q0.eps} \\end{tabular} \\end{center} \\caption{The energy loss rate versus $\\rho/\\mu_e$ due to pair annihilation (including radiative corrections), photoproduction and plasmon decay (solid lines) for several temperatures. The dotted lines refer to the analogous results of Ref. \\cite{Itoh96}, which also compute the rate for bremsstrahlung on nuclei. The effect of $\\Delta Q$ to pair annihilation can be appreciated in this logarithmic scale only for the largest temperature $10^{10} \\,^\\circ K$.} \\label{qcomp} \\end{figure} In order to perform a comprehensive study of order $\\alpha G_F^2$ contributions to the energy loss rate, we have also recalculated the contribution from $\\nu$-photo\\-production and plasmon decay processes. While the detailed analytical calculations will be discussed elsewhere \\cite{emmpp}, we here give our numerical results in Fig. \\ref{qcomp}, where we show the several contributions to $Q$. When possible, we also show for comparison the results of Ref. \\cite{Itoh96}. In particular our results for these processes agree quite nicely, at least in the whole region where they significantly contribute to the total energy loss rate. For the bremsstrahlung on nuclei we have used the analytic fitting formula of Ref. \\cite{Itoh96}. Fig. \\ref{qcomp} shows that, as well known, for large temperatures and not too high densities, pair annihilation dominates over the other two processes, while for low densities $\\nu$-photoproduction dominates over plasmon decay. On the other hand, for large densities the most relevant process is plasmon decay, whose rate however, along with those of all other processes, rapidly falls down for extremely high densities. This is a genuine plasma effect. Consider, for example, the behaviour of $\\nu$-photoproduction energy loss. As already noted in \\cite{Beaudet67b}, the decrease for very large densities is achieved only if one consistently takes into account the increasingly large photon thermal mass. In fact with a massless photon the $\\nu$-photoproduction curves in Fig. \\ref{qcomp} would rather reach a constant value. The main effect of $m_\\gamma^2$ is a lowering of the values of the Bose distribution function for photons, i.e. a smaller number of thermal photons. This reduces the energy loss rate induced by $\\nu$-photoproduction. \\begin{figure} \\begin{center} \\epsfysize=8.5cm \\epsfxsize=8.5cm \\epsffile{Qcontours.eps} \\caption{The regions in the $T-\\rho/\\mu_e$ plane where each of the processes i)-iv) contribute for more than $90 \\%$ to the total energy loss rate. We also show the contours for the relative correction $\\Delta Q/Q^0_{Tot}$ (see text) for the values $1\\%,0\\%,-1\\%,-2\\%,-3\\%,-4 \\%$.} \\label{qcontour} \\end{center} \\end{figure} In Fig. \\ref{qcontour} we show the regions in the temperature-density plane where a given process contributes to the total energy loss rate (including radiative corrections to pair annihilation) for more than 90\\%. We also summarize there our results on the radiative corrections to pair annihilation processes, by plotting the contours corresponding to $\\Delta Q/Q_{Tot}^0=1\\%,0\\%,-1\\%,-2\\%,-3\\%,-4 \\%$, where $Q_{Tot}^0$ is the total emission rate with pair annihilation calculated in Born approximation. These contours lie almost entirely in the region where the pair annihilation process gives the main contribution to $Q$. This result may affect the late stages of evolution of very massive stars by changing their configuration at the onset of Supernova explosion. This issue is presently under study." + } +} \ No newline at end of file